{
    "0206/astro-ph0206231_arXiv.txt": {
        "abstract": "{We continue the QSO search in the 10 square degrees Schmidt field around M\\,92 based on variability and proper motion (VPM) constraints. We have re-reduced 162 digitised $B$ plates with a time-baseline of more than three decades and have considerably improved both the photometric accuracy and the star-galaxy separation at $B>19$. QSO candidates are selected and marked with one out of three degrees of priority based on the statistical significance of their measured variability and zero proper motion. Spectroscopic follow-up observations of 84 new candidates with $B>19$ revealed an additional 37 QSOs and 7 Seyfert\\,1s. In particular, all 92 high-priority candidates are spectroscopically classified now; among them are 70 QSOs and 9 Seyfert\\,1s (success rate 86\\%). We expect that 87\\% (55\\%) of all QSOs with $B<19.0$ ($19.8$) are contained in this high-priority subsample. For the combined sample of high-priority and medium-priority objects, a completeness of 89\\% is estimated up to $B_{\\rm lim} = 19.5$. The sample of all AGNs detected in the framework of the VPM search in the M\\,92 field contains now 95 QSOs and 14 Seyfert\\,1s with $B\\le19.9$. Although the VPM QSOs were selected by completely different criteria, their properties do not significantly differ from those of QSOs found by more traditional optical survey techniques. In particular, the spectra and the optical broad band colours do not provide any hints on a substantial population of red QSOs up to the present survey limit. ",
        "introduction": "The number of known QSOs is growing very rapidly. Over the last decade, among others, the Durham/AAT survey (Boyle et al. \\cite{Boy90}), the Large Bright Quasar Survey (Hewett et al. \\cite{Hew95}), the Edinburgh Quasar Survey (Goldschmidt \\& Miller \\cite{Gol98}), and the Hamburg/ESO survey (Wisotzki et al. \\cite{Wis00}) have been completed. Presently, the 2dF Quasar Survey (Croom et al. \\cite{Cro01}) and the Sloan Digital Sky Survey (Schneider et al. \\cite{Sch01}) are extremely efficient at identifying very large numbers of quasars. The INT Wide Angle Survey (Sharp et al. \\cite{Sha01}) is expected to detect a statistically significant sample of high-redshift quasars. Very deep quasar samples were obtained in the Lockman hole via the X-ray satellite ROSAT (Hasinger et al. \\cite{Has98}) and in the optical domain with the Hubble Space Telescope (e.\\,g., Conti et al. \\cite{Con98}), respectively. Further, the VLA FIRST Bright Quasar Survey (e.\\,g., White et al. \\cite{Whi00}) will define a radio-selected QSO sample that is competitive in size with current optically selected samples. All these surveys select the QSO candidates on the basis of  their particular colours with respect to stars and galaxies or their brightness at X-ray or radio wavelengths. In other words, they rely on  the different broad-band spectral energy distribution (SED) as the prime selection criterion.  Despite the large number of QSOs now catalogued, the selection effects of the conventional surveys are not yet fully understood (see, e.g., Webster et al. \\cite{Web96}; White et al. \\cite{Whi00}; Gregg et al. \\cite{Gre02}). It is therefore important to perform alternative QSO surveys using selection criteria that do not directly rely on the shape of the SED. Variability and zero proper motion are such criteria. We started a variabillty and proper motion (VPM) survey based on several hundred digitised Schmidt plates (Meusinger et al. \\cite{Meu02}). Due to the special demands on the number and the time-baseline of the available observations such an attempt must be limited to comparatively small and confined areas.   We perform the VPM search in two fields of a size of ten square degrees each, centered on the globular clusters M~3 and M~92, respectively. The work in the M~92 field is the subject of the present series of papers. In Paper~1 (Brunzendorf \\& Meusinger \\cite{Bru01}), we described the observational material, the photometric and astrometric data reduction, and the selection procedure of the QSO candidates. The QSO sample resulting from the spectroscopic follow-up observations of these candidates was presented in Paper~2 (Meusinger \\& Brunzendorf \\cite{Meu01}). In Paper~3 (Meusinger \\& Brunzendorf \\cite{MeuB02}), the properties of the narrow-emission line galaxies among the VPM QSO candidates are discussed.  Both the completeness and the efficiency of the VPM search primarily depend on the photometric accuracy (see Paper~1). From the comparison with other optical QSO samples we find that the previous VPM QSO sample is virtually complete for brighter magnitudes, but its completeness drops rapidly at the faint end ($B>19$). In the context of the work on Paper~3, we have refined substantial parts of the reduction of all digitised plates with the result of a significantly improved sample of QSO candidates at fainter magnitudes. In the present paper, we present the results of the follow-up spectroscopy and discuss the properties of the enlarged VPM QSO sample.  In Sect.~2, we briefly describe the major revisions of the data reduction procedure. The new candidate selection and the selection effects are the subjects of Sect.~3. The new spectroscopic follow-up observations are described in Sect.~4. Section~5 deals with the newly detected QSOs and the properties of the enlarged QSO sample. Finally, conclusions are given in Sect.~6. As in the previous papers, we adopt $H_0 = 50$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ and $q_0=0$. ",
        "conclusions": "In our previous work (Papers~1 and 2), we have described a combined VPM QSO search on 162 digitised Schmidt plates of the M\\,92 field taken with the Tautenburg Schmidt in the $B$ band between 1960 and 1997. We have confirmed that this method provides an efficient technique for finding AGNs. The efficiency and the completeness of the VPM survey primarily depends on the photometric accuracy. The previous QSO sample suffered from substantial incompleteness at $B>19$. Therefore, we have revised the reduction of all $B$ plates on the basis of the SExtractor package (Bertin \\& Arnouts \\cite{Ber96}). Both the photometric accuracy and the star-galaxy separation at $B>19$ were substantially improved. The mean photometric accuracy is now $\\sigma \\approx 0.1$\\,mag at $B=19$ and $\\sigma \\approx 0.2$\\,mag at $B=20$. With the refined variability indices, a number of new QSO candidates of medium or high priority were selected. Spectroscopic follow-up observations of 84 new candidates revealed  an additional 37 QSOs and 7 Seyfert\\,1s with $B<19.9$.  The total VPM sample in the M\\,92 field now comprises 95 QSOs and 14 Seyfert\\,1s with $B\\le19.9$. For all these AGNs long-term lightcurves are available with a baseline of more than three decades. Among the 92 high priority QSO candidates with $18 \\le B \\le 19.8$,  we found 70 QSOs and 9 Seyferts\\,1s, corresponding to a high success rate of 86\\%. The completeness of the high-priority subsample is estimated to be about 87\\% for $B\\le19$ and 55\\% for $B\\le19.8$. A total completeness of 89\\% is derived for the sample of all QSOs brighter than $B=19.5$ from the comparison with the number counts given by  Hartwick \\& Schade (\\cite{Har90}). Remarkably, the properties of the VPM QSOs do not significantly differ from those of QSOs found by more traditional optical survey techniques. The deeper VPM sample from the present study confirms the result from Paper~2: the spectra and the optical broad band colour indices do not provide any hints on a substantial population of abnormal red QSOs up to the survey limit."
    },
    "0206/astro-ph0206141_arXiv.txt": {
        "abstract": "The problem of formation of the Ruderman-Sutherland type inner vacuum gap in neutron stars with ${\\bf\\Omega}\\cdot{\\bf B}<0$ is considered. It is argued by means of the condition $T_i/T_s>1$ (where $T_i$ is the critical temperature above which $^{56}_{26}$Fe ions will not be bound at the surface and $T_s$ is the actual temperature of the polar cap surface heated by the back-flow of relativistic electrons) that the inner vacuum gap can form, provided that the actual surface magnetic field is extremaly strong ($B_s\\gtrsim 10^{13}$~G) and curved (${\\cal R}<10^6$~cm), irrespective of the value of dipolar component measured from the pulsar spin down rate. Calculations are carried out for pulsars with drifting subpulses and/or periodic intensity modulations, in which the existence of the quasi steady vacuum gap discharging via ${\\bf E}\\times{\\bf B}$ drifting sparks is almost unavoidable. Using different pair-production mechanisms and different estimates of the cohesive energies of surface iron ions, we show that it is easier to form the vacuum gap controlled by the resonant inverse Compton scaterring seed photons than by the curvature radiation seed photons. ",
        "introduction": "The consecutive  subpulses in a sequence of single pulses of a number of pulsars change phase systematically between adjacent pulses, forming apparent driftbands of the duration from several to a few tenths of pulse periods. The subpulse intensity is also systematically modulated along driftbands, typically increasing towards the pulse centre. In some pulsars, which can be associated with the central cut of the line-of-sight trajectory, only the periodic intensity modulation is observed, without any systematic change of subpulse phase. On the other hand, the clear subpulse driftbands are found in pulsars associated with grazing line-of-sight trajectories. These characteristics strongly suggest an interpretation of this phenomenon as a system of subpulse-associated beams rotating slowly around the magnetic axis. \\citet[][ hereafter RS75]{rs75} proposed a natural explanation of the origin of subpulse drift, which involved a number of isolated ${\\bf E}\\times{\\bf B}$ drifting sparks discharging the quasi steady vacuum gap formed above the polar cap of the pulsar with ${\\bf\\Omega}\\cdot{\\bf B}<0$, in which $^{56}_{26}$Fe ions were strongly bound at the surface. Although the original idea of RS75 associating the rotating sub-beams with the circulating sparks is still regarded as the best model of drifting subpulse phenomenon, their vacuum gap was later demonstrated to suffer from the so-called binding energy problem \\citep[for review see][]{as91,um95,xqz99}. In fact, the cohesive energies of $^{56}_{26}$Fe ions used by RS75 proved largely overestimated and the inner vacuum gap envisioned by RS75 was impossible to form. However, it is worth emphasizing that RS75 considered the canonical surface dipolar magnetic fields with values determined from the pulsar spindown rate, although they implicitly assumed small radii of curvature ${\\cal R}\\sim 10^6$~cm required by the pair creation conditions, which is inconsistent with a purely dipolar field. Recently Gil \\& Mitra (2001; hereafter GM01) revisited the binding energy problem an found that the formation of a vacuum gap is in principle possible, although it requires an extremely strong non-dipolar surface magnetic field $B_s=bB_d$, where the coefficient $b \\gg 1$ in a typical pulsar, $B_d=6.4\\times 10^{19}(P\\dot{P})^{0.5}{\\rm G}=2\\times 10^{12}(P\\dot{P}_{-15})^{0.5}$~G is the dipolar field at the pole, $P$ is the pulsar period in seconds, $\\dot{P}$ is the period derivative and $\\dot{P}_{-15}=\\dot{P}/10^{-15}$.  In a superstrong surface magnetic field $B_s>0.1B_q$, where $B_q=4.414\\times 10^{13}$~G, the asymptotic approximation of \\citet{e66} used by RS75 in derivation of the height of quasi steady vacuum gap is no longer valid. In fact, in such strong field the high energy $E_f=\\hbar\\omega$ photons produce electron-positron pairs at or near the kinematic threshold $\\hbar\\omega=2mc^2/\\sin\\theta$, where $\\sin\\theta=h/{\\cal R}$, $h$ is the gap height, and ${\\cal R}={\\cal R}_610^6$~cm is the radius of curvature of surface magnetic field lines \\citep[e.g.][]{dh83}, $\\hbar$ is the Planck constant, $c$ is the speed of light, $m$ and $e$ are the electron mass and charge, respectively. The vacuum gap formed under such conditions was called the Near Threshold Vacuum Gap (hereafter NTVG) by GM01. They considered two kinds of high energy seed photons dominating the $e^-e^+$ pair production: the Curvature Radiation (CR) photons with energy $\\hbar\\omega=(3/2)\\hbar\\gamma^3c/{\\cal R}$ \\citep[RS75;][]{zq96}, and resonant Inverse Compton Scattering (ICS) photons with energy $\\hbar\\omega=2\\gamma\\hbar eB_s/mc$ \\citep{zq96,zetal97}, where $\\gamma$ is a typical Lorentz factor of particles within the gap. The corresponding vacuum gap is called the Curvature Radiation dominated (CR-NTVG) and the Inverse Compton Scattering dominated (ICS-NTVG), respectively. GM01 estimated the characteristic heights of both CR-NTVG and ICS-NTVG. In this paper we further refine these estimates by including the general relativistic (GR) effects of inertial frame dragging (IFD) and considering the heat flow conditions within the thin uppermost surface layer of the polar cap. Moreover, we use a broader range of cohesive energies of surface $^{56}_{26}$Fe ions. The obtained VG models are applied to pulsars with drifting subpulses and/or periodic intensity modulations, in which the presence of ${\\bf E}\\times {\\bf B}$ drifting spark discharges seems almost unavoidable (Deshpande \\& Rankin 1999, 2001; Vivekanand \\& Joshi 1999, and Gil \\& Sendyk 2000). ",
        "conclusions": "There is a growing evidence that the radio emission of pulsars with systematically drifting subpulses (grazing cuts of the line-of-sight) or periodic intensity modulations (central cuts of the line-of-sight) is based on the inner vacuum gap developed just above the polar cap \\citep{dr99,dr01,vj99,gs00}. To overcome the binding energy problem \\citet{xqz99,xzq01} put forward an attractive but exotic conjecture that pulsars showing the drifting subpulses represent bare polar cap strange stars (BPCSS) rather than neutron stars. However, as demonstrated in this paper, invoking the BPCSS conjecture is not necessary to explain the drifting subpulse phenomenon. The quasi steady vacuum gap, with either curvature radiation or inverse Compton scattering seed photons, can form in pulsars with ${\\bf\\Omega}\\cdot{\\bf B}<0$, provided that the actual surface magnetic field at the polar cap is extremely strong $B_s\\sim 10^{13}$~G and curved ${\\cal R}<10^6$~cm, irrespective of the value of dipolar component measured from the pulsar spindown rate. We have used two sets of the cohesive (bounding) energies of the surface iron ions: higher values obtained by AS91 and about five times lower values obtained by J86. If the actual cohesive energies are close to J86 values, then only ICS controlled VG can form and CR controlled VG never forms even under the most extreme conditions (see Fig. 1).  It is worth noting that ICS discussed in this paper is the so-called \"resonant\" scattering. The choice of this mode of ICS is justified by the superstrong magnetic fields relevant for the near-threshold regime. However, it is clear that the thermal photons also contribute some pair-producing gamma-rays \\citep[e.g.][]{zetal97}. This process cannot dominate the breakdown of the gap, since it requires higher surface temperatures, which may inhibit the formation of VG in the first place. One can only speculate that the dominating thermal ICS mode is associated with the pulse nulling phenomenon. These issues require further investigation.  The pulsars with drifting subpulses and/or periodic intensity modulation do not seem to occupy any particular region of the $P-\\dot{P}$ diagram (see Fig.~1 in GM01). Rather, they are spread uniformly all over the $P-\\dot{P}$ space, at least for typical pulsars (excluding young and millisecond pulsars, in which observations of single pulses are difficult in the first place). Therefore, it seems tempting to propose that radio emission of all pulsars should be driven by vacuum gap activities. An attractive property of such proposition is that the nonstationary sparking discharges induce the two-stream instabilities that develop at relatively low altitudes \\citep{u87,am98} where the pulsars radio emission is expected to originate \\citep{c78,kg97,kg98}. It is generally believed that the high frequency plasma waves generated by the two-stream instabilities can be converted into coherent electromagnetic radiation at pulsar radio wavelengths \\citep[e.g.][]{mgp00}. With such scenario, all radio pulsars would require a strong, non-dipolar surface magnetic field at their polar caps (e.g. Gil et al. 2002 a, b)."
    },
    "0206/astro-ph0206188_arXiv.txt": {
        "abstract": "Many black--hole sources emit a substantial fraction of their luminosities in blackbody--like spectral components. It is usual to assume that these are produced in regions at least comparable in size to the hole's Schwarzschild radius, so that a measure of the emitting area provides an estimate of the black hole mass $M$. However there is then no guarantee that the source luminosity (if isotropic) obeys the Eddington limit corresponding to $M$. We show that the apparent blackbody luminosity $L_{\\rm sph}$ and temperature $T$ must obey the inequality $L_{\\rm sph} < 2.3\\times 10^{44}(T/100\\ {\\rm eV})^{-4}$ ~erg~s$^{-1}$ for this to hold. Sources violating this limit include ultrasoft AGN and some of the ultraluminous X--ray sources (ULXs) observed in nearby galaxies. We discuss the possible consequences of this result, which imply either super--Eddington or anisotropic emission in both cases. We suggest that the ultrasoft AGN are the AGN analogues of the ULXs. ",
        "introduction": "Many luminous accreting sources have prominent spectral components which appear approximately blackbody. If the accretor is a black hole it is natural to assume that the dimensions of the emitting surface are comparable to the Schwarzschild radius, as this is the region where most of the gravitational potential energy is released. However the result of this procedure is not necessarily compatible with the Eddington limit on the source luminosity, derived from the requirement that the radiation pressure should not drive the accreting matter away. Examples of this for quasars have appeared in the literature (Puchnarewicz, 1994; Molthagen, Bade and Wendker 1998; Puchnarewicz 1998). Here we investigate the question systematically, and show that various classes of accreting black holes do not straightforwardly correspond to sub--Eddington blackbodies. In general the resolution of this difficulty appears to require super--Eddington or anisotropic emission. ",
        "conclusions": "Fig.~1 shows that the ultrasoft AGN have luminosities $L_{\\rm sph}$ a factor $\\sim 10 - 30$ above the value of $L_0$ appropriate to their measured temperatures, while several ULX sources show a similar if smaller effect. Before going further we should check if these results could be spurious. Probably the most serious possible cause of error arises from the claimed blackbody temperatures. These might be systematically too high either because (a) the values of $T$ are wrongly fitted, or (b) the spectra are not blackbody at all, but for example the result of Comptonization, or the effects of electron scattering in regions with low absorption opacity. Errors of the right order (factors $\\ga 3 - 10$) appear unlikely; and if the effect is a physical one we have to explain why it occurs only in a subset of AGN.  Nevertheless this area merits further study.  If we accept the values of $T$ appearing in Fig.~1, we should consider the possibility that the emission may come from a region of total area smaller than the Schwarzschild radius $R_s$, i.e. that $r < 1$. Since most of the accretion energy is released in a region of order $R_s$, this requires that it should be removed non--radiatively, and only converted to radiation in a much smaller region. The only way of ensuring this appears to involve magnetic fields. Merloni \\& Fabian (2001) have indeed suggested that magnetic reconnection in regions comparable with the local disc thickness $H$ may be the primary dissipation mechanism in accretion discs. However, they also require that these regions should be triggered at heights at least an order of magnitude larger than their size. Since $H \\sim 0.1R$, heights $\\ga R \\ga R_s$ are needed. In this picture, blackbody emission would result from reprocessing the primary emission from these `lamppost' regions on the disc surface. The reprocessing regions cannot be smaller than the lamppost heights $\\ga R_s$, implying $r > 1$ in this model. Similar considerations probably apply in other magnetic energy release pictures.  Thus assuming that the values of $T$ in Fig.~1 are not grossly in error, and that the radiating regions of the sources are unlikely to be significantly smaller than their Schwarzschild radii, we are left with the alternatives that the sources with $L_{\\rm sph} > L_0$ are genuinely super--Eddington, or radiate anisotropically. We note that either of these possibilities would probably remove the need for intermediate masses $M \\ga 100\\msun$ in even those ULXs which do not violate the limit $L_{\\rm sph} > L_0$. Begelman (2002) has proposed a mechanism allowing thin accretion discs to radiate at up to ten times the Eddington limit, while King et al. (2001) have suggested that most ULXs are anisotropic emitters.  The common feature here is that presumably both types of source are supplied with mass at rates close to or above the Eddington value $\\dot M_{\\rm Edd} \\simeq L_{\\rm Edd}/0.1c^2$. To zeroth order we would expect the resulting accretion geometry to be similar in the two cases despite the large difference in black hole mass. We therefore suggest that the ultrasoft AGN are the supermassive analogues of the ULXs."
    },
    "0206/astro-ph0206377_arXiv.txt": {
        "abstract": "{ We present the time integrated and time resolved spectral analysis of a sample of bright bursts selected with $F_{peak} \\ge 20\\ \\mathrm{phot\\ cm^{-2}\\ sec^{-1}}$ from the BATSE archive.  We fitted four different spectral models to the pulse time integrated and time resolved spectra. We compare the low energy slope of the fitted spectra with the prediction of the synchrotron theory [predicting photon spectra softer than $N(E)\\propto E^{-2/3}$], and test, through direct spectral fitting, the synchrotron shock model. We point out that differences in the parameters distribution can be ascribed to the different spectral shape of the models employed and that in most cases the spectrum can be described by a smoothly curved function. The synchrotron shock model does not give satisfactory fits to the time averaged and time resolved spectra. Finally, we derive that the synchrotron low energy limit is violated in a considerable number of spectra both during the rise and decay phase around the peak. ",
        "introduction": "The nature and emission mechanisms responsible for the prompt emission of Gamma--ray bursts (GRB) are still a matter of debate. From the phenomenological perspective much effort has been made in order to characterize and identify typical spectral properties of bursts, mostly by applying parametric but very general and simple spectral models. Probably the most widely adopted is that suggested by Band et al. (\\cite{Band b}), namely a smoothly connected double power law model. Its attractive feature is that it characterizes, within the observational energy window, the most relevant quantities, namely the peak energy, representing the energy at which most of the emission occurs, and the low and high energy components which are related, according to the most accredited emission theories, to the particle energy distribution and/or to the physical parameters of the emitting region. Indeed GRB spectra result to be well represented by the Band parameterization ($N(E)\\propto E^{\\alpha , \\beta}$), with typical low energy power law photon spectral indices $\\alpha$ between --1.25 and --0.25, high energy spectral indices $\\beta$ around --2.25 (Preece et al. \\cite{Preece1998}) and peak energy $E_{p}$ typically around 100--300 keV.  A fundamental  issue in the  spectral analysis of GRBs  is the integration timescale of  the spectra which are observed  to vary even on  millisecond timescales  (Fishman et  al. \\cite{Fishman  b}). BATSE spectra,  which  are integrated  for  a  minimum  of 128ms,  therefore represent  the best  way, presently available, to  constrain the  burst  emission mechanism. Furthermore,  in  order to  compare the  average spectral properties of different bursts, also time integrated spectra, covering the duration of the pulse or the entire burst, have been used in literature.  It has been found in general that the Band model not only well describes the time integrated spectra, but it also appears to fit the time resolved spectra of bright bursts. Alternative spectral models -- and in particular the predictions for synchrotron emission (Katz \\cite{Katz}, Crider at al. \\cite{Crider a},b) -- have also been recently tested on a large sample of time resolved spectra by Preece and collaborators (Preece et al. \\cite{Preece2000}).  Within  this scenario,  here  we  present the  study  of the  spectral properties  of  single pulses  within  bright  GRBs  which intends  to complement  the work mentioned  above, by  specifically aiming  at: 1) compare the  results of the  analysis of the  spectra averaged over  major pulses  in the  burst lightcurve  with the  time resolved spectra  of  the very  same  burst  in  order to  quantify  systematic differences; 2) consider both empirical (Band model) and more `physical' (synchrotron shock  model) spectral models for all spectra  and  compare  the  quality  of  the  corresponding  fits.  In particular  each spectrum  is fitted  with the  four models  we choose (Band's,  broken power--law,  thermal  Comptonization and  synchrotron shock model, the latter fitted  to temporal resolved BATSE spectra for the  first  time). We  also  examine  any  spectral `violation'  (with respect to  the predicted slope  in the case of  synchrotron emission) for the entire  burst evolution.  The development of  this work is the analysis and interpretation of  the spectral evolution morphologies of this sample of bursts.  The paper is structured as follows: In sect. 2 we describe the data and their selection criteria, while the spectral models adopted for the analysis are detailed in sect. 3. Section 4 presents the results of our work for the time integrated and time resolved spectra. Conclusions are drawn in sect. 5.  The time evolution of the spectral parameter will be the content presented in a following paper (Ghirlanda et al. in preparation). ",
        "conclusions": "We considered a sample of bright bursts detected by BATSE and performed a uniform analysis for the time integrated and the time resolved (typically 128 ms) spectra with four different models proposed in the literature.  We find that even with this time resolution no parametric model can better represent the data and different spectra require different shapes, re--confirming the erratic behaviour of bursts and also possibly indicating that time resolution on time--scales comparable with the variability one is needed to shed light on such erratic characteristics.  Indeed, an important result we confirm is that the average time integrated spectrum often used in the literature does not well represent the very same event resolved on shorter time-scales.  The time integrated spectra might still be used for a comparison of the average spectral shape among different pulses and as indicators of the average spectral parameters of the time resolved analysis although only the time resolved spectra should be used in any test of a physical emission model.  Finally, a considerable number of the fitted spectra are characterized by extremely hard low energy components with spectral index $\\alpha$ greater than $-2/3$, value predicted by synchrotron theory (Katz \\cite{Katz}). This violation was found by Crider et al. (\\cite{Crider a}) and has been recently reported by Frontera et al. (\\cite{Frontera}) in some GRBs observed by BeppoSAX. They report 1 sec time resolved spectra significantly harder than $E^{-2/3}$, mainly during the first phase of the burst emission. We have found that in 11 of the 25 bursts analyzed the $\\alpha$ limit violation is evident mainly in the spectra around the peak both during the rise and decay phase, and this could indicate that at least at some stages of the burst evolution -- possibly near the peak of emission itself -- radiative processes, other than synchrotron, can dominate the emission.  We also reported some examples of bursts which violate the synchrotron limit and are not characterized by the $\\alpha$--$E_{peak}$ anticorrelation predicted by the small pitch angle distribution synchrotron model proposed by Lloyd \\& Petrosian (\\cite{Lloyd}).  The obvious extension of this work, namely the study and discussion of the temporal evolution of the spectral shape, will be the subject of a forthcoming paper (Ghirlanda et al., in preparation)."
    },
    "0206/astro-ph0206296_arXiv.txt": {
        "abstract": "We present \\nnew\\ N$^0$ measurements from our HIRES/Keck database of damped \\lya abundances.  These data are combined with measurements from the recent and past literature to build an homogeneous, uniform set of observations. We examine photoionization diagnostics like Fe$^{++}$ and Ar$^0$ in the majority of the complete sample and assess the impact of ionization corrections on \\nalph\\ and \\alphH\\ values derived from observed ionic column densities of N$^0$, Si$^+$, H$^0$, and S$^+$.  Our final sample of \\nfin\\ \\nalph, \\alphH\\ pairs appears bimodal; the majority of systems show \\nalph\\ values consistent with metal-poor emission regions in the local universe but a small sub-sample exhibit significantly lower \\nalph\\ ratios. Contrary to previous studies of \\nalph\\ in the damped systems, our sample shows little scatter within each sub-sample. We consider various scenarios to explain the presence of the low \\nalph\\ sightlines and account for the apparent bimodality.  We favor a model where at least some galaxies undergo an initial burst of star formation with suppressed formation of intermediate-mass stars. We found a power-law IMF with slope 0.10 or a mass cut of $\\approx 5-8 \\msol$ would successfully reproduce the observed LN-DLA values. If the bimodal distribution is confirmed by a larger sample of measurements, this may present the first observational evidence for a top heavy initial mass function in some early stellar populations. ",
        "introduction": "\\label{sec:intro}  Abundance studies of the damped \\lya systems -- protogalaxies with large H\\,I surface densities $\\N{HI} > 2 \\sci{20} \\cm{-2}$ -- reveal the chemical enrichment history of the universe. These quasar absorption line systems dominate the neutral hydrogen reservoir to $z \\sim 4$ \\citep{wol95,storr00,peroux01} and, therefore, their metal abundances presumably reflect the past and current processes of protogalactic star formation \\citep[e.g.][]{pei99}.  High resolution surveys of the damped systems measure the metal content of these galaxies and trace the evolution in Zn and Fe metallicity with redshift \\citep{ptt94,pw00}.  Although these observations track the gross census of metals with time, they only crudely describe the physical processes of metal enrichment.  To examine issues related to the initial mass function (IMF), the formation epoch, and the star formation rate of individual systems, one must pursue other diagnostics.  One important avenue for addressing the details of chemical enrichment is through the investigation of relative metal abundances like O/Fe \\citep[e.g.][]{tinsley79}.  By comparing the so-called $\\alpha$-elements, elements presumed to form primarily in massive star supernovae, against Fe one roughly assesses the relative contribution of Type~Ia and Type~II SN which relates to the IMF and star formation history. Although current observations of the damped systems \\citep{pw02} suggest an $\\alpha$-enhancement at low metallicity similar to Galactic metal-poor stars \\citep{wheeler89,mcw97}, these measurements are subject to the effects of differential depletion \\citep{vladilo98,ledoux02}. To date, it has proved a great challenge to isolate the competing effects of nucleosynthesis and differential depletion for ratios like Si/Fe and O/Fe. One can minimize these uncertainties by examining a larger set of $\\alpha$ and Fe-peak elements \\citep[e.g.][]{dessauges02a}, but this requires extensive observations and is still subject to some uncertainty.  A complementary approach toward tracing the detailed enrichment history of protogalaxies is to consider the relative abundance of nitrogen. Nitrogen is mainly produced in the six steps of the CN branch of the CNO cycles within H burning stellar zones, where $^{12}$C serves as the reaction catalyst (see a textbook like Clayton 1983 or Cowley 1995 for a nucleosynthesis review).  Three reactions occur to transform $^{12}$C to $^{14}$N: $^{12}$C(p,$\\gamma$)$^{13}$N($\\beta$$^{+}$,$ \\nu$)$^{13}$C(p,$\\gamma$)$^{14}$N, while the next step, $^{14}$N(p,$\\gamma$)$^{15}$O, depletes nitrogen and has a relatively low cross-section. The final two reactions in the cycle transform $^{15}$O to $^{12}$C. Since the fourth reaction runs much slower than the others, the cycle achieves equilibrium only when $^{14}$N accumulates to high levels, and so one effect of the CN cycle is to convert $^{12}$C to $^{14}$N.  One issue in nitrogen evolution is to discover the source of the carbon which is converted into nitrogen, and of any oxygen which can contribute through the (slow) side chain $^{16}$O(p,$\\gamma$)$^{17}$F($\\beta$$^{+}$,$\\nu$)$^{17}$O(p,$\\alpha$)$^{14}$N. For example, stars may produce their own carbon (and some oxygen) during helium burning, and the carbon (and perhaps oxygen) is subsequently processed into $^{14}$N via the CN(O) cycle. In this case, nitrogen production is independent of the initial composition of the star in which it is synthesized and is referred to as {\\it primary} nitrogen. On the other hand, stars beyond the first generation in a galactic system already contain some carbon and oxygen, inherited from the interstellar medium out of which they formed. The amount of nitrogen formed from CNO cycling of this material will then be proportional to its C abundance (and also its O abundance, if the CNO cycling proceeds long enough to deplete the oxygen) and is known as {\\it secondary} nitrogen. To zeroth order, then, primary nitrogen production is independent of metallicity, while secondary production is linearly proportional to metallicity. The effects of metallicity on nitrogen production is one facet of nitrogen evolution we would like to understand clearly.  A second issue of concern in the origin of nitrogen is to identify the portion of the stellar mass spectrum which is most responsible for $^{14}$N production. While only massive stars (M$>$8M$_{\\sun}$) can produce the necessary internal temperatures required to synthesize most heavy elements, in the case of nitrogen intermediate-mass stars (IMS; 1-8M$_{\\sun}$) are also massive enough to attain required temperatures for its synthesis. In fact, both observational evidence and theoretical predictions strongly indicate that significant amounts of nitrogen are produced in intermediate-mass stars (M$<$8M$_{\\sun}$), while some is also synthesized in massive stars (Maeder 1992; Woosley \\& Weaver 1995; van~den~Hoek \\& Groenewegen 1997; Henry, Kwitter, \\& Bates 2000; Henry, Edmunds, \\& K{\\\"o}ppen 2000, hereafter HEK00; Marigo 2001; Meynet \\& Maeder 2002; Siess, Livio, \\& Lattanzio 2002). And since the timescale for IMS evolution is longer than for massive stars, nitrogen production and ejection into the interstellar medium may be significantly delayed with respect to oxygen and other heavy elements that are produced primarily in massive stars. This in turn means that nitrogen evolution is affected by the form of the initial mass function (IMF), time dependence of the star formation rate, and the effective stellar nitrogen yields of both massive and intermediate-mass stars.  Local measurements of emission line regions have examined the abundance of nitrogen relative to O and other $\\alpha$-elements as a function of metallicity (see Henry \\& Worthey 2000 for a compilation). These observations have provided evidence for the contributions of both the primary and secondary nitrogen mechanisms described above to the cosmic abundance evolution of this element. For example, one observes (i) a 'plateau' of \\nalph\\ measurements at [\\alphH]~$<-1$ consistent with primary nitrogen formation, i.e., independent of metallicity; and (ii) a rise in \\nalph\\ with increasing metallicity above [\\alphH]~$\\approx -1$ suggestive of secondary nitrogen formation. Similar observations can be performed with the damped \\lya systems (Pettini et al.\\ 1995; Lu et al.\\ 1998, hereafter L98; Centuri\\'on et al.\\ 1998). Because damped systems tend to have low metallicity ($Z < Z_\\odot/10$) their \\nalph\\ measurements probe the primary regime of nitrogen production. In addition, since these systems are relatively young and may not have achieved a stage of steady evolution because of differences in stellar evolution time scales over the stellar mass spectrum, N/$\\alpha$ may allow us to gauge the effects of time delay and the role of IMS in nitrogen production.  And unlike the $\\alpha$/Fe ratios described above, the \\nalph\\ observations are largely free of the uncertainties due to differential depletion because of the mild refractory nature of N, O, S, and Si. In fact the most significant source of uncertainty is due to the effects of photoionization: measurements of ions \\nti, \\oi, \\suii, H$^0$, and \\siii\\ must be converted to elemental abundances.  The initial studies of \\nalph\\ in damped systems suggested a significant scatter in the \\nalph\\ ratios (L98) with a few systems showing values below the plateau defined by low metallicity H~II regions. These observations are difficult to interpret as either purely primary or secondary nitrogen formation and a combination of the two channels was suggested (L98). In this paper, we will reexamine these conclusions. We present \\nnew\\ nitrogen measurements from our UCSD HIRES/Keck~I damped \\lya database (Prochaska et al.\\ 2001; hereafter, P01). and combine these data with \\nstud\\ from previous studies and an additional \\nlit\\ measurements from the recent literature.  We restrict our discussion to high resolution echelle observations because (1) the N\\,I transitions lie within the \\lya forest where line-blending is important; (2) several N\\,I transitions are subject to line saturation; and (3) the N\\,I 1134 and 1200 triplets are very closely spaced. We limit the analysis to the damped \\lya systems where the effects of photoionization are more likely to be small. We synthesize the abundance measurements from these various studies, investigate the effects of photoionization, and present a homogeneous, uniform sample of \\nalph\\ and \\alphH\\ measurements.  We then consider various models of N production and describe the implications for the star formation histories within the damped \\lya systems. ",
        "conclusions": "In this paper we presented \\nalph\\ vs.\\ \\alphH\\ measurements of a moderate sample of damped \\lya systems and identified a possible bimodality in the \\nalph\\ values at the lowest metallicity.  While the majority of low metallicity DLA have \\nalph\\ values consistent with low metallicity emission-line regions in the local universe, a sub-sample exhibits significantly lower \\nalph\\ values (the LN-DLA). Unlike previous works, we found minimal scatter among the members of each of the sub-samples and therefore argue that the DLA measurements form a bimodal distribution. We considered several scenarios which could account for these observations: (1) these systems are too young ($<250$~Myr) for their intermediate-mass stars to have produced N; (2) a reduced production of N in intermediate-mass stars at low metallicity; (3) Population III nucleosynthesis;  and (4) chemical evolution with a top-heavy IMF. Currently, we do not favor the first three scenarios.  If the LN-DLA are to be explained through the young age of the absorbing galaxy, then one would expect the LN-DLA to fill the parameter space beneath the plateau instead of comprising a bimodal distribution. Regarding the second option, it would be difficult to explain why DLA with the same metallicity exhibit such different \\nalph\\ values. Finally, the third avenue is promising but current models \\citep[e.g.][]{heger02} of Pop~III nucleosynthesis appear to underpredict the \\nalph\\ levels observed in the damped systems.  Therefore, we favor a scenario where the LN-DLA were enriched with an initial starburst with suppressed formation of intermediate-mass stars. We found that a starburst characterized by a power-law IMF with slope 0.10 or a mass cut of $\\approx 5-8 \\msol$ would successfully reproduce the observed LN-DLA values. This scenario allows for the bimodal DLA distribution of \\nalph\\ without requiring that we have observed the damped systems at a special moment in time. Furthermore, an initial burst of star formation in a low metallicity environment may suppress the formation of lower mass stars \\citep[e.g.][]{abel00,bromm01} and naturally lend to a bimodal distribution.  Before concluding, we wish to briefly comment on correlations between \\nalph\\ and other physical properties of the damped \\lya systems. In terms of kinematic characteristics, two of the LN-DLA systems (Q1946+75, z=2.8; Q2348--14, z=2.3) have among the simplest kinematics of any damped \\lya system. This may suggest these systems have particularly small mass and/or relatively quiescent star formation. We note, however, that several other systems with higher \\nalph\\ ratios also exhibit very simple kinematics. The only other trend we can identify is an increase in \\nalph\\ with increasing $z_{abs}$ which is contrary to expectation. Perhaps this absence of correlation is further evidence that the production of N is largely independent of environment and star formation history as suggested by the uniformity of the emission line sample."
    },
    "0206/astro-ph0206069_arXiv.txt": {
        "abstract": "We present polarimetric observations of 14 pre-main-sequence (PMS) binaries located in the Taurus, Auriga, and Orion star forming regions. The majority of the average observed polarizations are below 0.5\\%, and none are above 0.9\\%. After removal of estimates of the interstellar polarization, about half the binaries have an {\\it intrinsic} polarization above 0.5\\%, even though most of them do not present other evidences for the presence of circumstellar dust. Various tests reveal that 77\\% of the PMS binaries have or possibly have a variable polarization. LkCa~3, Par~1540, and Par~2494 present detectable periodic and phase-locked variations. The periodic polarimetric variations are noisier and of a lesser amplitude ($\\sim$0.1\\%) than for other types of binaries, such as hot stars. This could be due to stochastic events that produce deviations in the average polarization, a non-favorable geometry (circumbinary envelope), or the nature of the scatterers (dust grains are less efficient polarizers than electrons). Par~1540 is a Weak-line T~Tauri Star, but nonetheless has enough dust in its environment to produce detectable levels of polarization and variations. A fourth interesting case is W~134, which displays rapid changes in polarization that could be due to eclipses. We compare the observations with some of our numerical simulations, and also show that an analysis of the periodic polarimetric variations with the Brown, McLean, \\& Emslie (BME) formalism to find the orbital inclination is for the moment premature: non-periodic events introduce stochastic noise that partially masks the periodic low-amplitude variations and prevents the BME formalism from finding a reasonable estimate of the orbital inclination. ",
        "introduction": "Pre-main-sequence (PMS) stars are objects still contracting to the main sequence (MS) and are divided into 2 classes according to their masses: T~Tauri stars (TTS) are low-mass PMS stars with $0.5\\;$M$_{\\sun} \\lesssim\\;$M$\\lesssim 2.0\\;$M$_{\\sun}$, while Herbig~Ae/Be (HAeBe) stars represent their higher-mass counterparts with 2~M$_{\\sun} \\lesssim$ M $\\lesssim 10$ M$_{\\sun}$. These 2 classes of objects exhibit properties characteristic of their youth: (1) association with dark or bright nebulosities (remnant of the parent cloud), (2) emission line spectrum (spectral type F or later for the TTS, A or B for the HAeBe), (3) IR excesses, (4) position above the MS in the HR diagram, and (5) presence of the \\ion{Li}{1} absorption line. Both types of PMS stars show signs of circumstellar (CS) material in their environment, in the form of jets and bipolar flows, emission excesses in the IR, mm and sub-mm domains, P~Cygni or inverse P~Cygni profiles, linear polarization, and resolved disks or envelopes around some of the objects. In the case of binary stars, the disks can be found around each component (circumstellar (CS) disks) and/or around the binary system (circumbinary (CB) disks). Variability is observed in photometry, spectroscopy, and polarimetry.  TTS are further divided into 2 classes according to the width of their H$\\alpha$ emission line. Classical TTS (CTTS) have wide H$\\alpha$ emission lines with $W_{\\lambda}> 5$ \\AA\\, whereas the Weak-line TTS (WTTS) have $W_{\\lambda}<5$ \\AA\\ (Bertout 1989). Naked TTS (NTTS) are a subclass of the WTTS (Walter 1986; Wolk \\& Walter 1996), although some tend to use both terms as synonyms. Indeed, in addition to narrower H$\\alpha$ lines, the NTTS specifically do not show evidence for circumstellar (CS) material in their environments in the form of IR excesses, whereas the WTTS may or may not (Wolk \\& Walter 1996). WTTS and NTTS should not be confused with post-T~Tauri stars, which are still above the MS, but more evolved than TTS. For a review on T~Tauri stars and Herbig~Ae/Be stars, see for example Bertout (1989) and Catala (1989).  Dust grains produce polarization by scattering, polarization that has been known for a number of years (see Bastien 1996 for a review). In general, the red part of the visible spectrum of TTS exhibits linear polarizations of 1--2\\%, but the polarization distributions of CTTS and WTTS are markedly different: there are a few CTTS which have a very high polarization (more than 5\\%, sometimes up to 15\\%), but WTTS almost all have polarization levels below 2\\% (Bastien 1982, 1985; M\\'enard \\& Bastien 1992). Also, most TTS that have an active CS disk have higher polarization and near IR excess than other TTS (Yudin 2000). On average, HAeBe stars are more polarized than TTS (3.0\\% versus 1.6\\%), and there are clear differences between the polarization distributions for TTS and HAeBe stars (Yudin 2000). Most PMS stars have statistically higher polarization levels than more evolved stars which are closer to the MS and there is clear evidence of changes in polarimetric behavior of stars during the evolution from PMS to MS star; these changes in polarimetric behavior are related to the evolution of the CS environment (Yudin 2000).  The IR color indices and color excesses of TTS are well correlated with the polarization (Bastien 1985), implying that the same grains are responsible for both phenomenons. A compilation of polarization data for almost 500 TTS and HAeBe stars studied by Yudin (2000) reveals that for 85\\% of the sample stars, there is a correlation between the degree of polarization and the IR color index $(V - L)_{\\rm obs}$ and the color excess $E(V-L)$. This polarization is variable in a majority of objects: more than 85\\% of the TTS, and more than 70\\% of HAeBe stars present variability (Bastien 1988; M\\'enard \\& Bastien 1992). Variations are sometimes large and fast ($\\Delta P > 0.5$\\%, $\\Delta \\theta > 15\\arcdeg$, within 5 days, Bastien 1985), sometimes associated with luminosity and color changes (Grinin et al. 1991; Grinin, Kolotilov, \\& Rostopchina 1995), and in some cases periodic.  We have shown (Manset \\& Bastien 2000, 2001a, hereafter Paper~I and II) how the linear polarization of a binary surrounded by circumstellar matter varies periodically as a function of the orbital period, and how the geometry of the disks, the nature and characteristics of the scatterers, the masses of the stars, and the orbit characteristics affect the polarimetric curves. Models can be used to find the orbital inclination from observed polarimetric variations (see for example Rudy \\& Kemp 1978; Brown, McLean, \\& Emslie 1978). The work from Brown et al. (1978) (hereafter BME) uses first- and second-order Fourier analysis of the Stokes curves to give, in addition to the orbital inclination, moments related to the distribution of the scatterers in the CS and CB environments. The BME formalism was developed for Thomson scattering in optically thin envelopes, and for binaries in circular orbits. Since polarization in PMS stars is produced by scattering on dust grains, and most of the known spectroscopic PMS binaries have eccentric orbits, the BME formalism can not be used a priori. However, we have shown (Paper~I and II) that the BME analysis can still be applied in those cases, with a few limitations.  In this context, we have obtained polarimetric observations of 24 spectroscopic PMS binaries. A detailed analysis was presented for one of these binaries, the HAeBe star MWC~1080 (Manset \\& Bastien 2001b, hereafter Paper~III). Here we report the complete observations and detailed analysis for the PMS binaries located in the Taurus, Auriga, and Orion star forming regions (SFRs). Binaries located in the Scorpius and Ophiucus regions will be analyzed in a future paper. ",
        "conclusions": "We have presented polarimetric observations for 14 spectroscopic PMS binaries located in the Taurus, Auriga, and Orion SFRs. The majority of the PMS binaries observed have a detectable linear polarization; only LkCa~3, NTTS~045251+3016, Ori~569, and VSB~126 do not present polarization levels above the $3\\sigma$ limit. For most binaries, IS polarization is also an important component of the observations. After an estimation of this IS polarization is removed, a few (5 out of 14) of the binaries present clear indications of intrinsic polarizations $\\gtrsim 0.5\\%$ and significantly above the IS polarization: V826~Tau, UZ~Tau~E/W, DQ~Tau, GW~Ori, and Par~1540. Overall, 7 of our 14 binaries, including all the clearly identified CTTS binaries (UZ~Tau~E, DQ~Tau, and GW~Ori; Mathieu et al. 1997), show intrinsic polarizations above 0.5\\%. Interestingly, then, even the WTTS and NTTS, for which other types of observations are not indicative of significant amounts of circumstellar material, can have detectable levels of polarization.  As has been found for single PMS stars, the polarization of PMS binaries is variable. All binaries are statistically variable or suspected variable, except LkCa~3, NTTS~045251+3016, Par~2494, Ori~429, and VSB~126. Despite the results of those tests, graphs of polarimetric data as a function of orbital phase show that LkCa~3 could be variable in position angle, and Par~2494 is clearly variable in $\\theta$. Not enough data are available for DQ~Tau to determine its variability. When combining data from literature and our two observations for UZ~Tau~E/W, variability is present. Therefore, we find that 54\\% of the PMS binaries are variable (7 of out 13), or 77\\% (10 out of 13) variable or possibly so; these results are compatible with those of Bastien (1988) and M\\'enard \\& Bastien (1992). All the known CTTS binaries in our sample (except maybe DQ~Tau for which we only have one measurement) present polarimetric variations. Many NTTS or WTTS also show polarimetric variations. Therefore, around these stars, although it is not thought that there is much CS material, there is enough to produce polarimetric variations, as it can be shown that very little mass in the form of dust is needed to produce detectable levels of polarization.  Some stars have shown, once or even a few times, atypical values of polarization and/or position angle that are well below or above the rest of the data (Par~1540, Par~2494, and also MWC~1080, Paper~III). We believe these are real observations of events that strongly affected the stars and/or their environment. Single PMS stars are known sometimes to be strongly variable, so this is not a surprise. For the stars with a sufficient amount of data, additional observations at similar phases indicate that these atypical points are not related to the normal periodic behavior. Photometry or spectroscopy obtained at the same times as those of the atypical observations could help reveal the cause.  Only 6 binaries have enough observations to investigate the presence of periodic polarimetric variations: LkCa~3, V826~Tau, GW~Ori, Par~1540, Par~2494, and W~134. Statistical tests conclude that LkCa~3's and Par~2494's polarization is constant, but low amplitude periodic variations are nonetheless present. The position angle for LkCa~3's polarization outlines a sinusoidal wave between phases 0.2 and 0.65; since the data were obtained over 22 months, representing more than 600 orbits, it cannot be attributed to coincidence, but to a real variation. Although Par~2494's observations are rather noisy, binned data reveal low amplitude periodic variations in position angle. Therefore, low amplitude variations, especially if masked by noise, can be missed by statistical tests for variability, because these only consider the whole of the observations without their known order (as a function of phase).  A third binary, Par~1540, shows periodic variations, although they also are of low amplitude. Period analysis for Par~1540 and Par~2494, with PDM and LNP, confirm that the periodicity is due to the orbital motion, although not with a very high significance. Non-periodic or pseudo-periodic polarimetric variations could explain why it is difficult to see periodic variations and confirm them with period analysis methods. Those non-periodic variations may be caused by appearance or disappearance of hot or cool spots, flares, or major changes in the distribution or density of matter in the CS or CB environment. To avoid or at least decrease the effects of those events, observations should be carried on consecutive nights before any major stochastic change in polarization occurs, and until the orbital period has been sufficiently covered. Those three PMS binaries which exhibit low amplitude polarimetric variations are WTTS; this indicates that there is still enough dust in their environment to produce polarization and periodic variations.  A fourth interesting case is W~134, which may be an eclipsing system: two atypical position angles, taken more than a year apart, are seen near phase 0.55; such rapid changes are sometimes seen for eclipsing binaries (for example in the eclipsing binary EK~Cep, Manset \\& Bastien in preparation; see also St-Louis et al. 1993). The rapid change in position angle is not at the predicted phases for the eclipses; nonetheless, more polarimetric data should be taken near phase 0.55 to investigate the cause of the atypical polarimetric observations.  One of the goals of these observations was to find the orbital inclinations. Unfortunately, non-periodic or pseudo-periodic variations sometimes mask the truly periodic variations by introducing noise. This noise, as measured by 3 different methods, is too high for the BME formalism to find reasonable estimates of the orbital inclination. Three factors contribute to this difficulty. First, dust grains are the main scatterers in these systems, and it has been shown that dust grain produce polarimetric variations of smaller amplitude than electrons (Paper~II). Second, the disk around these short-period binaries are probably CB rather than CS ones, and CB disks produce variations of smaller amplitudes than CS disks (Paper~II). Finally, non-periodic events introduce noise that masks the already small amplitude variations. This last problem might only be improved by taking data on shorter periods of time.  Other parameters are returned by the BME formalism: $\\Omega$, the orientation of the orbital plane with respect to the plane of the sky, and moments of the distribution of the scatterers, used to measure the asymmetry with respect to the orbital plane ($\\tau_0 G$), and the degree of concentration towards the orbital plane ($\\tau_0 H$). Although the assumptions used in the BME formalism (scattering on electrons, circular orbits) are not met in the PMS binaries studied here and we have shown that even in simple cases the values are incorrect, the values returned for PMS binaries are of the same order of magnitude as values for other types of stars and as for our simulations.  A similar detailed analysis will be presented for PMS binaries in the Scorpion and Ophiucus SFRs in a coming paper."
    },
    "0206/astro-ph0206319_arXiv.txt": {
        "abstract": "The expansion of the Universe leads to a rapid drop in the hydrogen \\lya\\ effective optical depth of the intergalactic medium (IGM), $\\mtau\\propto (1+z)^{3.8}$, between redshifts 4 and 3. Measurements of the temperature evolution of the IGM and of the \\Hep\\ opacity both suggest that \\Hep\\ reionizes in this redshift range. We use hydrodynamical simulations to show that the observed temperature increase associated with \\Hep\\ reionization leads to a relatively sudden decrease in $\\mtau$ around the reionization epoch of $\\approx$ 10 per cent. We find clear evidence for such a feature in the evolution of $\\mtau$ determined from a sample of $\\sim 1100$ quasars obtained from the SDSS. \\Hep\\ reionization starts at redshift $\\approx 3.4$, and lasts for $\\Delta z \\approx 0.4$. The increase in the IGM temperature also explains the widths of hydrogen absorption lines as measured in high-resolution spectra. ",
        "introduction": "Neutral hydrogen in the intergalactic medium (IGM) resonantly scatters the flux blueward of the hydrogen \\lya emission line in quasar spectra (Gunn \\& Peterson 1965; Bahcall \\& Salpeter 1965; Peebles 1993, \\S 23). The fact that not all the flux is absorbed implies that the IGM is very highly ionized. The neutral fraction is determined from the balance between photo-ionizations, produced by the UV-background radiation from galaxies and quasars, and recombinations. For hydrogen, the recombination rate depends on temperature $\\propto 1/T^{0.7}$, and so an increase in temperature will lead to a decrease in the level of absorption. Recent evidence suggests that \\Hep\\ reionizes around a redshift $z\\sim$3--3.5, and the associated temperature increase should have a measurable effect on the mean absorption. In this letter, we compute the optical depth evolution of a simulation in which \\Hep\\ reionization heats the IGM at redshift 3.4. The resulting dip in the evolution of the mean absorption matches that recently measured by Bernardi et al.\\ (2002) in the Sloan Digital Sky Survey (SDSS) data.  Hydrodynamical simulations have proved to be very successful in reproducing many properties of the observed absorption (Cen et al. 1994; Zhang, Anninos \\& Norman 1995; Miralda-Escud\\'e et al.\\ 1996; M\\\"ucket et al.\\ 1996; Hernquist et al.\\ 1996; Theuns et al.\\ 1998; Machacek et al.\\ 2000; see e.g. Efstathiou, Schaye and Theuns (2000) for a recent review). Most of the absorption at redshifts 2--5 arises in modestly over- and under dense filamentary and sheetlike structures, leading to a forest of \\lya-absorption lines in the spectra of quasars (Lynds 1971). These structures can be described and understood with simple physical models in which the typical size of the absorbers is determined by the Jeans length of the photo-heated gas (Bi \\& Davidsen 1997; Schaye 2001). Given that these hydrodynamical simulations reproduce the data in great detail, we are confident that they will also allow us to investigate the effect of a sudden increase in $T$ on the mean absorption.  Several lines of recent evidence suggest that \\Hep\\ reionizes around $z\\sim $ 3--3.5, significantly later than hydrogen. Observations of the \\Hep\\ \\lya-forest detect a sudden increase in the mean \\Hep\\ opacity around $z\\sim 3$ (Reimers et al.\\ 1996; Heap et al.\\ 2000; Kriss et al. 2001; Smette et al.\\ 2002). Such a rapid increase in the flux of \\Hep\\ ionizing photons is expected to occur at the final stages of reionization when individual \\Hepp\\ bubbles around sources percolate (Gnedin 2000). The resulting hardening of the ionizing background could be responsible for the observed jump in the relative abundance of C{\\sc IV}/Si{\\sc IV} (Songaila \\& Cowie 1996), although more recent data (Kim et al.\\ 2002) apparently do not seem to support such a change (see also Dav\\'e et al.\\ 1998). Finally, entropy injection associated with \\Hep\\ reionization will increase the temperature of the IGM (Miralda-Escud\\'e \\& Rees 1994), which has the effect of making the \\lya-absorption lines broader on average. The gas will become nearly isothermal when the change in temperature is large. Schaye et al.\\ (2000) detected both signatures in high-resolution Keck data, using a method based on the finding of Schaye et al.\\ (1999; see also Ricotti et al.\\ 2000 and Bryan \\& Machacek 2000) that the variation of the widths of the absorption lines as a function of column density, can be used to measure the temperature as a function of the density. Theuns et al.\\ (2002a) performed a wavelet analysis of the spectra of several high resolution QSO spectra, and found that the data required a large increase in $T$ by a factor $\\sim 2$ around $z\\sim 3.3$, consistent with late \\Hep\\ reionization (Miralda-Escud\\'e \\& Rees 1994). The temperature decrease after $z\\sim 3.3$ is also consistent with this interpretation (Theuns et al.\\ 2002b). Other heating mechanisms have been discussed in the literature as well, for example galactic winds (Cen \\& Bryan 2001). Although these may have contributed to the entropy of the high redshift gas, it seems unlikely they could result in a sudden change of temperature as seems required by the data.  In the next section, we use hydrodynamical simulations to investigate the effect of \\Hep\\ reionization on the evolution of the mean effective optical depth, \\mtau\\ . Section 3 summarizes the method used by Bernardi et al.\\ (2002) to measure $\\mtau(z)$ from a sample of 1061 SDSS quasar spectra, and section 4 compares the data to the simulations. ",
        "conclusions": ""
    },
    "0206/astro-ph0206363_arXiv.txt": {
        "abstract": "s{The spin-up of a neutron star crucially depends on the maximum orbital frequency around it, as do a host of other high energy accretion phenomena in low mass X-ray binaries, including quasi-periodic oscillations (QPOs) in the X-ray flux.  We compare the maximum orbital frequencies  for MIT-bag quark stars and for neutron stars modeled with the FPS equation of state. The results are based on relativistic calculations of constant baryon sequences of uniformly rotating strange star models, and are presented as a function of the stellar rotational frequency. The marginally stable orbit is present outside quark stars for a wide range of parameters, but outside the FPS neutron stars it is present only for the highest values of mass. This allows a discrimination between quark stars and neutron stars in the resonance theory of kHz QPOs.} ",
        "introduction": "If the recently discovered kHz QPOs~\\cite{Klis00} are a manifestation of strong gravity, they may be used to constrain the external metric of the compact source and the equation of state of matter at supranuclear densities.  Klu{\\'z}niak et al.\\cite{KluznMW90} suggested that the mass of a neutron star may be derived by observing the maximum orbital frequency, if it occurs in the marginally stable orbit, and that the low frequency QPOs occurring in X-ray pulsars will have their counterpart in LMXBs, at frequencies in the kHz range. Such kHz QPOs have indeed been discovered and their frequency used to derive mass values of about $2M_\\odot$, under the stated assumption.\\cite{Kaaret97,Zhang98,Kluzn98,Bulik00} In fact, the QPO frequency may correspond to a larger orbit, and hence a smaller mass.\\cite{KluznA02}  The question whether quark stars may have maximum orbital frequencies as low as the observed kHz QPO frequencies has also been investigated. Bulik et al.\\cite{BGK99a,BGK99b} showed that slowly rotating strange stars, described by the simple MIT bag model with massless and non-interacting quarks, have orbital frequencies at the marginally stable orbit higher than the maximum frequency of 1.07  kHz in 4U 1820-30, reported by Zhang et al.\\cite{Zhang98}.  However, the ISCO frequencies can be as low as 1 kHz when more sophisticated models of quark matter (with massive strange quarks and lowest order QCD interactions) and/or rapid stellar rotation are taken into account.\\cite{StergKB99,GondeSBKG01,ZduniBKHG00,ZdunHGG00} The  lowest orbital frequency at the ISCO was found to be attained either for non-rotating massive configurations close to their  maximum mass limit, or for configurations at the equatorial mass-shedding limit for a broad range of stellar masses. The effect of the crust, if present, has been investigated for normal evolutionary sequences.\\cite{ZdunHG01}  Typically the crust increases the maximum orbital frequency at the Keplerian limit.   If the strange star configurations described by Dey et. al.\\cite{Dey98} are allowed, the rotational and maximal orbital frequencies are much higher~\\cite{Datta00,GondeBZGRDD00,GondeBKZG01} than those for neutron star models, or for  the MIT-bag models of quark stars. The maximum orbital frequencies for the Dey et al. models are always higher than the kHz QPO frequencies observed to date, and higher than 1.5 kHz for  stars with masses greater than $1 \\, M_\\odot$.    \\begin{figure} \\vskip 2.5cm \\psfig{figure=procMRSS.eps,height=2.5in}\\qquad \\psfig{figure=procMRNS.eps,height=2.45in} \\caption{Gravitational mass versus radius for sequences of constant baryon mass.  Each sequence is labeled by this baryon mass in solar mass units, as well as (in parentheses) by the gravitational mass of the static configuration, if it exists. The angular momentum increases along each sequence from $J=0$ for static configurations, or $J_ {\\rm min}$ for supramassive stars, to $J_{\\rm max}$ for the mass-shedding limit (thick solid line).  The thick dashed lines indicate stars marginally stable to axisymmetric perturbations.  The left panel shows quark stars modeled with the MIT-bag e.o.s. for massless non-interacting quarks. The right panel: neutron stars modeled with the FPS e.o.s. The short-dashed lines correspond to rotating models for which the marginally stable orbit does not exist. The dotted lines separate the models with and without a marginally stable orbit.} \\label{fig:MR} \\end{figure}  \\begin{figure} \\vskip 2.5cm \\psfig{figure=procfmsfSS.eps,height=2.5in}\\qquad \\psfig{figure=procfmsfNS.eps,height=2.5in} \\caption{Maximum orbital frequency vs. the frequency of rotation $\\Omega/2\\pi$ for the sequences of Fig.~1.  One sequence for a very low mass quark star (with $M_{\\rm b}=0.01 M_\\odot$) is shown, the critical point on this sequence for Newtonian dynamical instability to non-axisymmetric perturbations is indicated by an asterisk and dynamically unstable configurations are denoted with the dotted line. The various dashed lines have the same meaning as in Fig.~1} \\label{fig:fmsf} \\end{figure} ",
        "conclusions": ""
    },
    "0206/astro-ph0206155_arXiv.txt": {
        "abstract": "We develop the formalism necessary to study four-point functions of the cosmic microwave background (CMB) temperature and polarization fields.  We determine the general form of CMB trispectra, with the constraints imposed by the assumption of statistical isotropy of the CMB fields, and derive expressions for their estimators, as well as their Gaussian noise properties.  We apply these techniques to initial non-Gaussianity of a form motivated by inflationary models.  Due to the large number of four-point configurations, the sensitivity of the trispectra to initial non-Gaussianity approaches that of the temperature bispectrum at high multipole moment.  These trispectra techniques will also be useful in the study of secondary anisotropies induced for example by the gravitational lensing of the CMB by the large scale structure of the universe. ",
        "introduction": "}  Beyond the power spectra of the cosmic microwave background (CMB) lies the relatively unexplored territory of non-Gaussian statistics.  Studies of its non-Gaussianity hold the potential to reveal physics at the two ends of time.  Non-Gaussianity in the primary anisotropies from the recombination epoch can test the inflationary model of the origin of fluctuations (e.g., \\cite{SandB1,Gangui1,Gangui2,Falk1,Munshi1}). Non-Gaussianity in the secondary anisotropies, arising during the transit of a CMB photon through the large-scale structure of the universe, probes the nature of the dark energy and dark matter (e.g., \\cite{Bernardeau1,PandC,Spergel2,Zaldarriaga1,HuDE}).  The primary challenge facing non-Gaussian studies of the CMB is the selection of an appropriate statistics.  The term ``non-Gaussianity'' tells us what the distribution is not, not what it is. Like the power spectra, the higher-point correlations of the multipole moments of CMB fields provide a set of statistics with definitive predictions in the cases of cosmological interest.  Unlike the power spectra, there are a large number of potential observables, associated with the distinct configurations of the points, requiring the development of new techniques for their prediction and estimation.  In particular, it is important to identify the symmetry properties of the spectra to build optimal statistics for the detection of non-Gaussianity.  Non-Gaussian signatures in the three-point function or bispectrum of the temperature distribution \\cite{Gangui1,Spergel2,Verde1, Cooray1,Komatsu1} and polarization \\cite{Hu2} as well as techniques for their extraction \\cite{Luo1,Heavens1,Spergel1,Sandvik1,Phillips1,Komatsu3} have been studied extensively in the past few years. The four-point function or trispectrum has recently received much attention due to its use in the study of the gravitational lensing of the CMB \\cite{Bernardeau1,Zaldarriaga1,HuTrispec,Cooray2}. Techniques for measuring certain components have been tested on the Cosmic Background Explorer (COBE) data \\cite{KomatsuPhD,Kunetal01}. Still, a complete treatment incorporating the full symmetry properties of the temperature and polarization fields has been lacking in the literature.  In this paper we complete the formalism established for the temperature trispectrum \\cite{HuTrispec}.  The addition of polarization information leads to a multiplicity of trispectra corresponding to all possible combinations of three observable fields.   It has been recently shown that the higher point correlations of the CMB polarization contain the majority of the information on gravitational lensing in the CMB \\cite{HuOka02}. Trispectra also quantify the non-Gaussian errors to temperature and polarization power spectra measurements.  The outline of the paper is as follows. We consider the general symmetry and noise properties of trispectra in Sec.~\\ref{Sect:Formalism}. As an illustration of the construction and noise properties of trispectra, we apply these techniques to the initial non-Gaussianity in the curvature fluctuations of the form predicted by slow-roll inflation in Sec.~\\ref{Sect:Initial}. We show that the sensitivity to initial non-Gaussianity in the trispectra can approach that in the temperature bispectrum \\cite{Komatsu1}. In Appendix \\ref{Appendix:SpinS} we summarize relations useful for the study of high order correlations in the polarization.  In Appendix \\ref{Appendix:Additional} we cover the details in the properties, measurement, and approximation of the trispectra that may be useful for future studies. ",
        "conclusions": "}  We have introduced a complete formalism for the study of 4-point correlations in the CMB temperature and polarization fields.  This formalism should be useful in future tests of the non-Gaussianity of the CMB induced in the early universe and by the evolution of structure. It is also of use in determining the non-Gaussian contributions to errors in temperature-polarization power spectra measurements.  We have applied these techniques to a particular form of trispectra motivated by inflation, generalizing previous treatments to higher order in the initial nonlinearity of the curvature fluctuations. Typical slow-roll inflationary models predict an amplitude to the trispectra that is far from observable and so a detection of this type of non-Gaussianity would rule out a large class of models.  We have shown that because of the large number of trispectra configurations, the sensitivity to initial non-Gaussianity in the trispectra approach that of the well-studied temperature bispectrum at high multipoles.  Trispectra from secondary anisotropies such as gravitational lensing \\cite{HuTrispec}  are expected to be substantially larger and should be fruitful ground for future studies.   While measurement of these non-Gaussian signatures will no doubt prove challenging due to foregrounds, systematic effects and computational cost, the wealth of information potentially contained therein may justify the large effort that will be required."
    },
    "0206/astro-ph0206225_arXiv.txt": {
        "abstract": "We describe the algorithm that selects the main sample of galaxies for spectroscopy in the Sloan Digital Sky Survey from the photometric data obtained by the imaging survey.  Galaxy photometric properties are measured using the Petrosian magnitude system, which measures flux in apertures determined by the shape of the surface brightness profile. The metric aperture used is essentially independent of cosmological surface brightness dimming, foreground extinction, sky brightness, and the galaxy central surface brightness.  The main galaxy sample consists of galaxies with $r$-band Petrosian magnitude $r \\leq 17.77$ and $r$-band Petrosian half-light surface brightness $\\muh \\leq 24.5$ magnitudes per square arcsec.  These cuts select about 90 galaxy targets per square degree, with a median redshift of 0.104.  We carry out a number of tests to show that (a) our star-galaxy separation criterion is effective at eliminating nearly all stellar contamination while removing almost no genuine galaxies, (b) the fraction of galaxies eliminated by our surface brightness cut is very small ($\\sim 0.1\\%)$, (c) the completeness of the sample is high, exceeding 99\\%, and (d) the reproducibility of target selection based on repeated imaging scans is consistent with the expected random photometric errors.  The main cause of incompleteness is blending with saturated stars, which becomes more significant for brighter, larger galaxies. The SDSS spectra are of high enough signal-to-noise ratio ($S/N > 4$ per pixel) that essentially all targeted galaxies $(99.9\\%)$ yield a reliable redshift (i.e., with statistical error $< 30\\rm \\, km\\,s^{-1}$). About 6\\% of galaxies that satisfy the selection criteria are not observed because they have a companion closer than the $55''$ minimum separation of spectroscopic fibers, but these galaxies can be accounted for in statistical analyses of clustering or galaxy properties.  The uniformity and completeness of the galaxy sample make it ideal for studies of large scale structure and the characteristics of the galaxy population in the local universe. ",
        "introduction": "The Sloan Digital Sky Survey (SDSS; \\citealt{York}) is carrying out an imaging survey in five photometric bands of $\\pi$ ster in the north Galactic cap, and a follow-up spectroscopic survey of roughly $10^{6}$ galaxies and $10^{5}$ quasars, complete within precisely defined selection criteria. The main scientific drivers of the SDSS are the large-scale distributions of galaxies and quasars.  In order to carry out precise measurements of galaxy clustering on the largest scales, and to measure the distribution of galaxy properties with the highest possible precision, it is necessary that the sample of galaxies for which spectra are taken be selected in a uniform and objective manner. The northern spectroscopic survey targets two samples of galaxies: a flux-limited sample to $r=17.77$ (hereafter called the main sample) and a flux- and color-selected sample extending to $r=19.5$, designed to target luminous red galaxies (LRGs).   This paper describes the algorithm used to select the main galaxy sample and presents demonstrations that the algorithm meets the survey goals of uniformity and completeness.  A separate paper (\\citealt{LRG}) discusses the LRG sample.  \\subsection {The Sloan Digital Sky Survey} \\label{sec:SDSS}  The SDSS hardware, software, and data products are summarized by \\citet{York} and \\citet{EDR}. In brief, the survey is carried out using a dedicated, wide-field 2.5m telescope, a mosaic CCD camera (\\citealt{Gunn}), two fiber-fed double spectrographs, and an auxiliary 0.5 m telescope for photometric calibration. The imaging is done in drift scan mode with the 30 photometric CCDs of the mosaic camera imaging $\\approx 20$ square degrees per hour, in five broad bands, $u,g,r,i$ and $z$ \\citep{Fukugita:1996} that cover the entire optical range from the atmospheric ultraviolet cutoff in the blue to the sensitivity limit of silicon CCDs in the red. The imaging data are 95\\% complete for point sources at $r^* \\approx 22.2$, and the photometric calibration is accurate to about 3\\% in $r$ at this writing (\\citealt{Hogg:2001,Smith:2002}). Because this calibration is still preliminary, we will refer to current measurements with the notation $u^*$, $g^*$, $r^*$, $i^*$, $z^*$, but we use $u$, $g$, $r$, $i$, $z$ to refer to the SDSS filter and magnitude system itself.\\footnote{This notation is a change from some earlier papers, including \\citet{Fukugita:1996}, which referred to the filter system as $u^\\prime$, $g^\\prime$, $r^\\prime$, $i^\\prime$, $z^\\prime$; see the discussion by \\citet{EDR}.} The astrometric calibration \\citep{Pier:2002} is done by comparison with the Tycho-2 \\citep{hoeg00} and UCAC \\citep{zacharias00} standards, and is accurate to 0.1 arcsec rms per coordinate.  The imaging data are reduced using a series of interlocking pipelines (\\photo; \\citealt{adass}), which flat-field the data, find all objects, match up detections in the different bands and perform measurements of their properties, and apply the photometric and astrometric calibrations. Spectroscopic targets --- the galaxies described in this paper, LRGs (\\citealt{LRG}), quasars (\\citealt{QSO}), and a variety of other categories of objects \\citep{EDR}, are chosen from the resulting catalog of detected objects.  The spectroscopic component of the survey is carried out using two fiber-fed double spectrographs, covering the wavelength range 3800\\AA\\ to 9200\\AA\\ over 4098 pixels.  They have a resolution $\\lambda/\\Delta\\lambda$ varying between 1850 and 2200, and together they are fed by 640 fibers, each with an entrance diameter of $3''$.  The fibers are manually plugged into plates inserted into the focal plane; the mapping of fibers to plates is carried out by a tiling algorithm (\\citealt{tiling}) that optimizes observing efficiency in the presence of large-scale structure.  The finite diameter of the fiber cladding prevents fibers on any given plate from being placed closer than $55''$ apart.  For any given plate, a series of fifteen-minute exposures is carried out until the mean signal to noise ratio (S/N) per resolution element exceeds 4 for objects with fiber magnitudes (i.e., as measured through the $3''$ aperture of the fiber) brighter than $g^*=20.2$ and $i^*=19.9$, as determined by preliminary reductions done at the observing site. Under good conditions (dark, clear skies and good seeing), this typically requires a total of 45 minutes of exposure.  \\subsection{The Main Galaxy Spectroscopic Sample}  The main galaxy spectroscopic survey is fully sampled to its magnitude limit within the survey footprint, which is planned to be an elliptical area of extent $110^\\circ\\times 130^\\circ$, chosen to minimize Galactic extinction and maximize observing efficiency. The median redshift of this sample is $z \\approx 0.1$.  This large galaxy sample will allow us to measure many independent modes of the density fluctuations on scales comparable to the peak of the galaxy power spectrum, largely free from the aliasing that can affect surveys with at least one narrow dimension (cf., Kaiser \\& Peacock 1991; Tegmark 1995).  For some instrumental set-ups and scientific goals (e.g., low-order measures of large scale clustering), one can gain efficiency by sparse sampling, \\ie, by observing only a fraction of the galaxies down to some limiting magnitude (cf., Kaiser 1986). However, sparse sampling adversely affects other kinds of investigation, including group and cluster studies, high-order clustering measures, and recovery of the underlying galaxy density field (see, e.g., \\citealt{SS96}).  Moreover, the field of view and number of spectroscopic fibers of the SDSS were chosen to allow simultaneous spectroscopy of essentially {\\em all} the galaxies in a given field to the faintest magnitude for which the 2.5m telescope can measure redshifts in a reasonable amount of time.  We have therefore opted for complete sampling in the main galaxy redshift survey.  We wish to select a magnitude-limited galaxy sample.   We have carried out the selection in a single observed band for simplicity. We wish the galaxy detection and photometric measurement in that band to be of high S/N, and we prefer a red passband so that $K$ corrections are modest, fluxes are determined mainly by the older stars that dominate the stellar mass, and uncertainties in Galactic reddening make little difference to the inferred galaxy magnitude. In the SDSS filter system, this implies either the $r$ or $i$ band. We adopt the former because the sky background is brighter and more variable in the $i$ band than in the $r$ band. The use of a red bandpass tilts the sample slightly towards galaxies of earlier morphological type, but at these bright magnitudes, the $g-r$ color distribution of galaxies is quite narrow \\citep{Ivezic02}, and the distribution of galaxy types is not radically different from what we would obtain with $g$-band selection (see also Fig.~4 of \\citealt{Shimasaku01}).  Although we will detail a number of subtleties below, the basic procedure that we use to define galaxy magnitudes and select spectroscopic targets can be summarized as follows. Star-galaxy separation is carried out by comparing the exponential or de Vaucouleurs model magnitude of an object to its Point Spread Function (PSF) magnitude.  We define the (angular) Petrosian radius $\\RP$ of a galaxy to be the radius at which the local surface brightness in an annulus about $\\RP$ is 1/5 of the mean surface brightness within $\\RP$.  We define the $r$-band Petrosian magnitude of a galaxy, $r_P$, based on the flux within a circular aperture of radius $2\\,\\RP$.  In the absence of seeing effects, the Petrosian magnitude measures the light within a well-defined metric aperture on any given galaxy which is independent of its redshift or foreground extinction.  We define the half-light surface brightness $\\muh$ to be the mean surface brightness within a circular aperture containing half of the Petrosian flux.  The main galaxy sample consists of galaxies with $r_P \\leq 17.77$ and $\\muh \\leq 24.5$ magnitudes per square arcsec, after correcting for Galactic extinction following Schlegel, Finkbeiner, \\& Davis (1998; hereafter SFD).  The outline of this paper is as follows.  In \\S~\\ref{sec:properties}, we describe our goals for the target selection algorithm. We discuss the measurements of Petrosian quantities in detail in \\S~\\ref{sec:petrosian} and Appendix~\\ref{sec:petro-appendix}.  The target selection algorithm itself is described in \\S~\\ref{sec:algorithm}. Various tests to show that the algorithm meets the survey requirements are described in \\S~\\ref{sec:test}. We conclude in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}   \\subsection{Summary of the algorithm performance}  The main spectroscopic galaxy sample of the SDSS is a reddening-correct $r$-band magnitude limited sample of galaxies brighter than $r_{P} = 17.77$, with an estimated surface density of 92 galaxies per square degree. The magnitude is measured within a Petrosian aperture, so as to provide a meaningful measure of a fraction of the total light of the galaxy that is independent of distance to the galaxy, reddening, and sky background. Star-galaxy separation is based on the difference between PSF and galaxy model magnitudes, which effectively quantifies the extension of the source relative to a PSF.  We reject objects with Petrosian half-light surface brightness $\\muh > 24.5$, a cut that eliminates $\\sim 0.1\\%$ of galaxies brighter than the magnitude limit. In the range $23 < \\muh < 24.5$, we use a measure of the difference between local and global sky brightness to increase our efficiency of targeting real galaxies.   We have objectively tested the star-galaxy separation algorithm and the completeness and reproducibility of the spectroscopic sample using imaging and spectroscopic data taken during the commissioning phase of the survey.  During commissioning, we refined the criteria in the target selection algorithm to achieve our goals on the completeness and efficiency of the spectroscopic sample. At the time of this writing, we find that the star-galaxy separation is accurate to better than $2\\%$, with the main contaminants being close double stars. The fraction of true galaxies rejected by the star-galaxy separation criterion is only $\\sim 0.3\\%$.  The completeness of the main galaxy sample is a function of magnitude. At bright magnitudes ($r^* < 15$), we find that we target $95\\%$ of the galaxies in the Zwicky catalog, while the remaining $5\\%$ are missed because they are blended with saturated stars. From comparison with visual inspection of bright galaxies ($r^* < 16$) over 200 square degrees of sky, we find that the  completeness increases to about $97.6\\%$. Finally, from comparison with a visual inspection of all objects brighter than $r^* = 18$ over 22 square degrees of sky, we find that the completeness of the galaxy sample to the magnitude limit is above $99\\%$. The only significant source of incompleteness that we have identified is blending with saturated stars; this incompleteness is higher for brighter galaxies because they subtend more sky.  Essentially all main sample galaxies (99.9\\%) that are observed spectroscopically yield successful redshifts.  About 10\\% of galaxy targets do not receive a fiber on the first spectroscopic pass because they lie within $55''$ of another sample galaxy. Some of these galaxies lie in regions of plate overlap and are observed subsequently, and the fraction of galaxies that are missed in the end because of the fiber separation constraint is about 6\\%. These missing galaxies can be accounted for in any statistical analysis by appropriate weighting of the galaxies in close pairs that are observed.  We have tested the reproducibility of the galaxy sample by selecting targets from repeated scans of the same region of the sky. We find that $94.5\\%$ of the spectroscopic sample galaxies are selected in both the scans. About $3.7\\%$ of galaxies fall out of the sample because they cross the magnitude limit and are replaced by a similar number of galaxies crossing in the other direction; this fraction is consistent with expectations based on random errors in the Petrosian magnitudes. Other galaxies fall out of the sample because of changes in saturation or star-galaxy separation.  Reproducibility of target selection is therefore high, and the random photometric errors that lead to non-reproducibility are not expected to cause systematic biases in statistical analyses.   \\subsection{Scientific applications of the SDSS imaging and spectroscopic data}  The imaging data on which we tested and refined the galaxy target selection algorithm, and the resulting galaxy spectroscopic sample have been studied in the context of both large scale structure and properties of galaxies. Extensive tests by Scranton \\etal\\ (2001) show that the imaging data obtained by the SDSS are free from internal and external systematic effects that influence angular clustering for galaxies brighter than $r^* = 22$, almost four magnitudes below the limit of the spectroscopic sample.  At the bright end, \\citet{Yasuda} studied the bright galaxy sample in the same data, and showed that the photometric pipeline correctly identifies and deblends blended objects and provides correct photometry for bright $(r^* < 16)$ galaxies.  The spectroscopic galaxy sample targeted using development versions of the target selection algorithm during the commissioning phase of the survey has been used to measure the luminosity function of galaxies as a function of surface brightness, color, and morphology (Blanton \\etal\\ 2001). A primary goal of the SDSS is to measure the properties of large scale structure as traced by different types of galaxies. Zehavi \\etal\\ (2002) used this spectroscopic sample to measure the correlation function and pairwise velocity dispersion of samples defined by luminosity, color, and morphology. Bernardi \\etal\\ (2002) used the spectra and photometry to study the correlations of elliptical galaxy observables including the luminosity, effective radius, surface brightness, color, and velocity dispersion. All these studies show that the galaxies targeted spectroscopically by the SDSS constitute a uniformly selected sample spanning a wide range of galaxy types, ideal for analyses of large scale structure and galaxy properties."
    },
    "0206/astro-ph0206013_arXiv.txt": {
        "abstract": "{ RX~J1914.4+2456 and RX~J0806.3+1527 have been proposed as double degenerate binaries with orbital periods of 569s and 321s respectively. An alternative model, in which the periods are related to the spin of a magnetic white dwarf in an intermediate polar system, has been rejected by other authors. We show that a face-on, stream-fed intermediate polar model for the two systems is viable and preferable to the other models. In each case, the X-ray modulation periods then represent the rotation of the white dwarf in the binary reference frame. The model explains the fully modulated X-ray pulse profiles, the X-ray spectra, the antiphase between X-ray and optical/infrared modulation, the lack of longer period modulation, and the low level of polarization. The optical spectrum of RX~J0806.3+1527 suggests that Balmer series lines are present, blended with HeII lines. This is unlike the spectra of any of the known AM CVn stars and suggests that the system is not a double degenerate binary. The optical spectrum of RX~J1914.4+2456 has spectral features that are consistent with those of a K star, ruling out the double degenerate models in this case. The lack of optical/infrared emission lines in RX~J1914.4+2456 may be attributed to a high mass accretion rate and its face-on orientation. Its reported period decrease may be a short term spin-up episode driven by the current high $\\dot{M}$. Finally we suggest that there is an observational selection effect such that the face-on intermediate polars that are detected will all have a stream-fed component, and the purely stream-fed intermediate polars that are detected will all be face-on systems. ",
        "introduction": "The X-ray sources RX~J1914.4+2456 (hereafter RXJ1914, but now also known as V407 Vul) and RX~J0806.3+1527 (hereafter RXJ0806) each display a single coherent X-ray and optical modulation with a period of order several minutes and no other confirmed modulations in any waveband observed. The most widely accepted models interpret the periods as due to orbital motion in a double degenerate binary system.  In this scenario, the two systems have the shortest orbital periods of any known binary star. A review of both systems is presented by Cropper, Ramsay \\& Wu (2003).  In the rest of this section we summarise the observational history of the two objects and note the difficulties in reconciling their behaviour with that of double degenerate binary systems. In section 2 we present simple analytical estimates of the physical properties of face-on stream-fed intermediate polars (IPs) and simulations demonstrating the accretion flow in these systems. We then consider how the various observational characteristics of RXJ1914 and RXJ0806 may be understood in terms of this model.  \\subsection{RX J1914.4+2456}  RXJ1914 was identified from the {\\em ROSAT} all sky survey as a 569s X-ray pulsator and suggested initially to be a member of the then recently recognised class of soft IPs, with 569s representing the spin period of a white dwarf (Motch et al 1996).  Subsequent {\\em ROSAT} observations failed to reveal any longer periods, apparently ruling out the IP model because an orbital period modulation might be expected. The unusual X-ray pulse profile, with zero flux for half the cycle, was claimed to require a system close to $90^{\\circ}$ inclination in conflict with the lack of X-ray eclipses; a model in which the 569s period is the beat period was also ruled out (Cropper et al 1998). Instead, a double degenerate polar model was suggested, with a 569s orbital period -- the first magnetic analogue of the AM CVn stars (Cropper et al 1998). Optical and infrared spectroscopy and photometry of the optical counterpart revealed that the V-band through to J-band modulations are all roughly anti-phased with the X-ray modulation, that the optical and infrared spectra show no emission lines, and that the system exhibits negligible polarisation (Ramsay et al 2000; 2002a). As this cast doubt on the double degenerate polar model, alternative interpretations were suggested including a double degenerate Algol (direct impact accretor) system (Ramsay et al 2002a; Marsh and Steeghs 2002) and a double degenerate electrically powered system (Wu et al 2002). Further doubt regarding the double degenerate accretor models (polar or Algol type) was raised by the discovery that the X-ray modulation frequency of RXJ1914 is increasing. If this period decrease represents a secular evolutionary trend, the observation is in direct conflict with a system which accretes via Roche lobe overflow from a degenerate star, since in that case a secular orbital period increase would be expected (Strohmayer 2002). Very recently, Strohmayer (2004) has claimed that a power spectrum of the {\\em Chandra} X-ray data shows evidence for a sideband structure to the 569s signal. This indicates a previously unseen longer period in the system of around $\\sim 1$ hour. The data also confirm the steady decrease of the 569~s period, with a frequency derivative of $6 \\times 10^{-18}$~Hz~s$^{-1}$. A final piece of evidence that poses a question for all three double degenerate models was provided by a spectrum obtained by Danny Steeghs (private communication) which has spectral features that are consistent with those of a K star.  \\subsection{RX J0806.3+1527}  RXJ0806 was discovered as a 321s X-ray pulsator amongst serendipitous X-ray sources observed by the {\\em ROSAT} HRI, and suggested as an IP (Israel et al 1999). Its X-ray modulation is remarkably similar to that of RXJ1914, showing a 50\\% duty cycle with the flux reduced to zero between pulses and it too was suggested to be a double degenerate polar (Burwitz \\& Reinsch 2001; Israel et al 2002).  The faint optical counterpart (Israel et al 1999; Burwitz \\& Reinsch 2001) displays an optical period coincident with that seen in X-rays and no convincing longer period (Ramsay, Hakala \\& Cropper 2002b, Israel et al 2002), although Reinsch (2003) reports some evidence for a possible 4700~s period in both optical and X-ray data that was also hinted at in the earlier observations.  As with RXJ1914, RXJ0806 displays antiphased X-ray and optical pulse profiles (Israel et al 2003a). Unlike RXJ1914, the spectrum of RXJ0806 shows faint, broad emission lines superimposed on a blue thermal continuum (Israel et al 2002). Further optical spectroscopy by Reinsch (2003) rules out a main sequence secondary of any spectral type, but allows the possibility that the system could contain a brown dwarf donor star.  Polarimetric data obtained by Israel et al (2003b) show that RXJ0806 exhibits no circular polarization, but that linear polarization is present at a level of 1.7\\%$\\pm$0.3\\%. It has been reported (Strohmayer 2003; Hakala et al 2003) that RXJ0806 shows a decreasing period, which if interpreted as a secular change in the orbital period would again rule out a system in which material is accreted from a degenerate companion (double degenerate polar or Algol models). However Woudt \\& Warner (2003) cast doubt on the measurement of this period derivative claiming that the period count is too uncertain. ",
        "conclusions": "It is difficult to see how the double degenerate models are viable given the observation of hydrogen lines in the spectrum of RXJ0806 (Israel et al 2002) and the spectral features consistent with those of a K star seen in RXJ1914 (Steeghs, private communication). If the period decreases seen in RXJ1914 (Strohmayer 2002) and RXJ0806 (Stroh\\-meyer 2003; Hakala et al 2003) are confirmed, and do indeed represent secular changes in the period, this too would rule out accretion from a degenerate donor star and the identification of the periods as orbital in nature. We note though that such period changes may also be understood in terms of magnetic cycles on the donor star, so this evidence is not compelling. The lack of polarization in RXJ1914 and low level polarization in RXJ0806 are a problem for the double degenerate polar model, whilst the geometrical constraints argue against the double degenerate Algol model.  For both RXJ1914 and RXJ0806, the X-ray pulse profiles, X-ray spectra, lack of other modulation periods, antiphased optical/infrared modulation, and level of polarization may be understood in terms of a face-on, stream-fed intermediate polar model. The implied orbital periods in each case are below the period gap (RXJ1914: $P_{\\rm orb} < 2.37$~h; RXJ0806: $P_{\\rm orb} < 1.33$~h), and the mass ratios ($M_2/M_1$) are relatively high. In the case of RXJ1914, the lack of spectral lines and the observed rate of period decrease may all be attributed to a current phase of high mass accretion rate, with $\\dot{M} \\sim 10^{17}$~g~s$^{-1}$. The possible detection of sideband structure in the X-ray power spectrum of RXJ1914 (Strohmayer 2004) and the possible long period X-ray modulation in RXJ0806 (Reinsch 2003) each suggest that the originally detected short period is related to a white dwarf spin period rather than an orbital period. Finally, the preferential detection of stream-fed systems with a low inclination angle may be the result of observational selection effects.  If the face-on stream-fed IP model is correct, then the emission lines seen in RXJ0806, and any which are ever detected from RXJ1914, should exhibit no sinusoidal radial velocity variations at the 321s or 569s periods, but may show modulation at a longer orbital period. If, instead, 321s and 569s do represent orbital periods in double degenerate binaries then sinusoidal radial velocity variations should be present at these periods.  In the face-on, stream-fed IP model, there may be a kinematic signature of the flipping accretion stream in any lines which are found. For half the beat cycle, the stream flows essentially towards the observer before turning to crash onto the upper, facing, magnetic pole of the white dwarf. For the other half of the beat cycle, the stream flows away from the observer, before turning to crash onto the lower, hidden, magnetic pole.  This stream-flipping may lead to velocity variations of spectral lines. Any variation would be at the 569s or 321s period in the two systems but would not be a simple sinusoidal modulation."
    },
    "0206/gr-qc0206054_arXiv.txt": {
        "abstract": "Quantum geometry predicts that a universe evolves through an inflationary phase at small volume before exiting gracefully into a standard Friedmann phase. This does not require the introduction of additional matter fields with ad hoc potentials; rather, it occurs because of a quantum gravity modification of the kinetic part of ordinary matter Hamiltonians. An application of the same mechanism can explain why the present-day cosmological acceleration is so tiny. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206081_arXiv.txt": {
        "abstract": "We report on high-resolution spatial and timing results for binary millisecond pulsar PSR J0218$+$4232 obtained with the Chandra HRC-I and HRC-S in imaging mode. The sub-arcsecond resolution image of the HRC-I (0.08-10 keV) showed that the X-ray emission from \\psr is consistent with that of a point source excluding the presence of a compact nebula with a size of about $14\\arcsec$ for which we had indications in ROSAT HRI data. The presence of a DC component is confirmed. This X-ray DC component has a softer spectrum than the pulsed emission and can be explained by emission from a heated polar cap. With our HRC-S observation we obtained a 0.08-10 keV pulse profile with high statistics showing the well-known double peaked morphology in more detail than before: The two pulses have broad wings and pulsed emission appears only to be absent in a narrow phase window of width $\\ltap 0.1$. The absolute timing accuracy of $\\sim 200 \\mu$s makes it possible to compare for the first time in absolute phase the X-ray pulse profile with the highly structured radio profile and the high-energy $\\gamma$-ray profile (0.1-1 GeV). The two X-ray pulses are aligned within absolute timing uncertainties with two of the three radio pulses. Furthermore, the two $\\gamma$-ray pulses are aligned with the two non-thermal X-ray pulses, corresponding to a probability for a random occurrence of 4.9$\\sigma$, strengthening the credibility of the earlier reported first detection of pulsed high-energy $\\gamma$-ray emission from a (this) millisecond pulsar. ",
        "introduction": "\\begin{figure*}[t] \\hbox{\\hspace{0.2cm}{\\psfig{file=idl_psrj0218+4232_hrci_map_lin_5x5_bw.ps,width=8.0cm,height=8.0cm}} \\hspace{1.00cm} {\\psfig{file=idl_psrj0218+4232_hrci_radprf_25_0.5.ps,width=8.0cm,height=7.5cm}} }  \\caption{(left) Chandra 0.08-10 keV HRC-I image of a $10\\arcsec\\times 10\\arcsec$ region centered on the radio pulsar position of \\psr. The radio position is marked with a `+' sign. The angular distance between the radio pulsar position and the X-ray centroid is $\\sim 0\\farcs 6$, consistent with the Chandra localization accuracy. (right) Radial distribution of HRC-I (Level 1) events using the best fit maximum likelihood X-ray position as centre. Superposed as dotted line is the radial profile from the PSF determined from the 2D fit procedure, which is consistent with the Chandra response to a point source.  The dashed line indicates the background level derived from counts in the range 10 -- 25 arcsec from the centre. \\label{hrcispatial} } \\end{figure*}   \\psr is a 2.3 ms pulsar in a two day orbit around a low mass ($\\sim 0.2$ M$_{\\sun}$) white dwarf companion \\citep{nav1995,kerk1997}. Pulsed X-ray emission with a Crab-like double pulse profile has been reported from ROSAT 0.1-2.4 keV data \\citep{kuip1998a} and BeppoSAX MECS 1.6-10 keV data \\citep{min2000}. The pulsed spectrum as measured by the MECS appeared to be remarkably hard with a power-law photon index $0.61 \\pm 0.32$, harder than measured for any other radio pulsar. Furthermore, \\citet{kuip2000} report the detection with EGRET of pulsed high-energy (0.1-1 GeV) $\\gamma$-ray emission from this ms pulsar. Their argument is based on three lines of evidence: (1) the 0.1-1 GeV data show a 3.5$\\sigma$ pulsation at the radio period; (2) the $\\gamma$-ray light curve resembles the one seen in hard X-rays, namely a phase separation of $\\sim$ 0.45 between two pulses/maxima; (3) the spatial analysis shows that the position of the EGRET source 3EG J0222+4253 moves from the position of the nearby BL Lac 3C 66A towards the pulsar position with decreasing gamma-ray energy (for energies between 100 and 300 MeV all source counts could be attributed to PSR J0218+4232). They also showed that the two $\\gamma$-ray pulses/maxima appeared to be aligned in absolute phase with two of the three radio pulses detected at 610 MHz. Unfortunately, the timing accuracies of the ROSAT and BeppoSAX observations were insufficient to construct X-ray profiles in absolute phase. None of the current models for pulsed X-ray and $\\gamma$-ray emission from radio pulsars offers a consistent explanation for the above summarized high-energy results on \\psr \\citep{kuip2000}.  \\psr is also remarkable in that it is the only Crab-like ms pulsar with a large DC (unpulsed) fraction of $63 \\pm 13\\%$ in the ROSAT band below 2.4 keV \\citep{kuip1998a}, as well as a large DC fraction of $\\sim 50\\%$ in radio, systematically over the range 100-1400 MHz \\citep{nav1995}. The DC components as measured in the ROSAT and radio observations could be explained by emission from a compact nebula with diameter $\\sim 14\\arcsec$, but in both cases the indications were at the limit of the imaging capabilities. Assuming that the radio DC component is compact, combined with the measured very broad and structured radio pulse profile, \\citet{nav1995} suggested that the magnetic field of \\psr is almost aligned with the rotation axis, the observer viewing the system under a small angle with respect to the rotation axis. \\citet{stairs1999} measured the magnetic inclination angle analyzing radio polarization profiles. Their rotation vector model fits indicate that the magnetic inclination angle is indeed consistent with $0\\arcdeg$ ($8\\arcdeg\\pm11\\arcdeg$). Unfortunately, in their fits the line-of-sight inclination is unconstrained. If the DC component in X-rays is also compact, for the suggested geometry of a nearly aligned rotator and a small viewing angle, it can originate in the pulsar magnetosphere as well as from a heated polar cap of the neutron star.  The objectives of our Chandra observations were: 1) To establish the spatial extent of the X-ray DC component, compact or extended; 2) To construct an X-ray pulse profile which can be compared in absolute phase with radio profiles and the 3.5$\\sigma$ EGRET high-energy (0.1-1 GeV) $\\gamma$-ray profile. ",
        "conclusions": "Our first objective of the Chandra observations was to determine the spatial extent of the X-ray DC component and to find a likely explanation for its origin. As is shown in Fig. 1, no evidence for extended emission near \\psr at $\\sim 1\\arcsec$ scales (diameter) is seen, rejecting the indication for a compact nebula found in our analysis of ROSAT HRI data \\citep{kuip1998a}. The combined ROSAT, BeppoSAX MECS and Chandra data led us to conclude that the DC spectrum is significantly softer than the pulsed spectrum, the latter being clearly of non-thermal origin. If we adopt the geometry proposed by \\citet{nav1995}, a nearly aligned rotator and a small viewing angle, the soft DC X-ray component can be explained as thermal emission from the polar cap of the neutron star which stays visible to the observer all the time. For old millisecond pulsars the thermal emission can only originate from reheating of the polar cap area by back-flowing accelerated particles. In fact, for a ms-pulsar the polar-cap half angle $\\theta_{pc}$ ($=\\arcsin(\\sqrt(R_{ns}\\cdot 2\\pi/p\\cdot c))$) extends over a relatively large angle, e.g. $17\\fdg 5$ for \\psr. This means that the spin-axis is well within the (magnetic) polar-cap region, for the measured angle between the spin- and magnetic axes amounts $\\sim 10\\arcdeg$ ($8\\arcdeg\\pm 11\\arcdeg$) \\citep{stairs1999}. Therefore, a small observer viewing angle to the system (angle spin-axis and line of sight) can keep the observer looking at the heated polar region continuously.  The \\citet{nav1995} radio observations had a VLA beam size of $16\\arcsec$, also allowing an interpretation of the radio DC component (reported DC fraction $\\sim 50\\%$) as a compact nebula up to the size of the VLA beam. No new observations have been reported with a smaller beam size, but \\citet{stairs1999} and \\citet{kuzmin2001} revisited the pulsed emission. We fitted the total radio spectrum of \\citet{nav1995} (five data points) with a power-law shape, index $2.57 \\pm 0.07$, and the new total pulsed radio spectrum (three data points from \\citet{nav1995}, two from \\citet{stairs1999} and one from \\citet{kuzmin2001}), index $2.58\\pm 0.15$. The spectra appeared identical in shape and the updated DC fraction is $(13 \\pm 9)\\%$, significantly lower than the earlier estimate. This leads us to the conclusion that the radio DC component is most likely also compact and has the same magnetospheric origin as the pulsed emission, which has a very structured and remarkably broad profile and is indeed practically never ``off''.  Our second objective was to obtain with Chandra a more significant X-ray profile than measured sofar, to study the profile structure in more detail in the X-ray band below 10 keV, and to have for the first time absolute timing, allowing multiwavelengths phase comparisons. The profile in Fig. \\ref{hrcspulseprof} with a significance of $15.2\\sigma$ is indeed much more significant than the MECS profile ($6.8\\sigma$). Now we can really see that the two pulses have broad wings and that the profile reaches the DC level only in a very narrow phase interval around phase 0.35 (phase extent $\\ltap 0.1$). In fact, Fig.\\ref{hepulsprofstack} shows that the radio and the X-ray profiles both reach a minimum level for approximately the same {\\it absolute} phase region. Furthermore, Fig.\\ref{hepulsprofstack} shows that the X-ray pulses are aligned in absolute phase with two of the three radio pulses and the two $\\gamma$-ray pulses within the timing uncertainties of the different measurements.  \\psr is the first and only millisecond pulsar for which we reported evidence for detection of high-energy gamma-ray emission up to 1 GeV \\citep{kuip2000}. We take the alignment in absolute phase of the non-thermal X-ray and $\\gamma$-ray pulses as important supporting evidence for our first detection of high-energy $\\gamma$-rays from a millisecond pulsar. In Sect. 5 we showed that the probability that a random timing signal reaches a $3.5\\sigma$ modulation significance {\\em and} has its 55 pulsed excess counts in phase with the two X-ray pulses amounts $\\sim 4.9\\sigma$.  \\citet{wall2000} reported the detection of short-term variability in the high-energy $\\gamma$-ray emission from 3EG J0222+4253 (3C66A/PSR J0218$+$4232) in a systematic search for all 170 unidentified sources in the 3rd EGRET Catalog \\citep{hart1999}. They noted that if this flaring is due to \\psr: ``It would be an unusual source in two ways: it would be the only millisecond pulsar seen by EGRET and the only pulsar to show strong flaring\". We argue, however, that in their search the reported evidence for flaring is not significant: For each of the 170 unidentified sources \\citet{wall2000} produced light curves per viewing period (VP) with 2 day flux values each time they were in the field-of-view during a VP. For 3EG J0222+4253 (3C66A/PSR J0218$+$4232) they found the most significant evidence for variability, namely in VP15 with a variability index V=2.6, corresponding to a $\\sim 3\\sigma$ significance for a random detection in a {\\em single} trial assuming Gaussian statistics (although Poisson statistics apply). This variability was due to one single high 2 day flux value. However, this source was viewed in 4 observations (no indication for variability in the other VPs) which makes the probability ($4\\times 0.002512$) 1\\% or 2.58$\\sigma$ to find V=2.6 in 1 out of 4 observations. Furthermore, \\citet{wall2000} analysed 144 VPs with duration larger than 3 days producing for the 170 unidentified sources few-hundred light curves, making the probability to find once V=2.6 (or one single high 2 day flux value) few times unity! They correctly noted that their Monte Carlo probabilities give misleading values (too optimistic) as the used averages having no uncertainties assigned. Also in the latter case the very large number of ``trials'' has been ignored. Since there was no {\\em a priori} reason to select 3EG J0222+4253 for a single trial, we donot regard the indication for short-term variability of this source significant.  The X-ray results on the DC emission and the radio results on the DC and pulsed emission suggest for \\psr an emission scenario of a nearly aligned rotator and a small viewing angle. The latter angle can still be as large as $\\sim 20\\arcdeg$ to explain the X-ray DC emission if this originates from a polar-cap heated by particle bombardment. For explaining the pulse profile with two hard-spectrum X-ray/$\\gamma$-ray peaks it is more critical to know this viewing angle to get a handle on the geometry. For both competing classes of models, polar cap models (PC; e.g. \\citet{daugherty94,daugherty96}) and outer gap models (OG; see e.g. \\citet{cheng86a,cheng86b,ho89}) production of hard X-ray/$\\gamma$-ray emission in the magnetospheres of millisecond pulsars has been predicted (e.g. \\citet{bhattacharya91,sturner94}). Recently, millisecond pulsars among them \\psr were considered for PC models by \\citet{dyks99,bulik99,bulik00,zhang00}. For \\psr they did not succeed in reproducing the measured high-energy spectral shape. More recently, \\citet{dyks02} succeeded in reproducing the high-energy spectrum of PSR J0218$+$4232, but required non-orthodox assumptions about the electron energy distribution or emission altitude as well as off-beam viewing geometry. \\citet{wozna02} reproduced the double-peak profile with the measured phase separation for a small inclination angle of $8\\arcdeg$ and a viewing angle of $29\\arcdeg$ (angles approximately consistent with our findings). However, for these angles the spectrum reached its maximum luminosity for too high energies around 100 GeV. The fact that in the case of PC models a double-peak hard X-ray/$\\gamma$-ray profile can be modelled for a nearly aligned rotator with a $\\sim 20\\arcdeg$ viewing angle follows already from \\citet{daugherty96}, who showed that the $\\gamma$-ray beam produced at the polar cap rim is relatively wide, amounting $\\sim 26 \\arcdeg$ ($1.5 \\times \\theta_{pc}$) for PSR J0218$+$4232. For emission produced at higher altitudes in the magnetosphere (as is proposed in the more recent versions of PC models), the beam will become even wider. In all these discussions the real viewing angle in combination with the inclination angle determines whether the measured double peak profile can be obtained. Also for OG models a double peak profile is in principle possible, depending on the actual viewing angle and which part of the outer gap (in altitude) is visible to the observer. For OG models the hard X-ray/$\\gamma$-ray cone might become too broad for the case of \\psr as a nearly aligned rotator if the hard X-ray/$\\gamma$-ray production takes place only close to the light cylinder. However, \\citet{hirotani02} argued recently that the $\\gamma$-ray production is not limited to regions in the outer gap above the null-charge surface, but can also originate from altitudes closer to the neutron star. More detailed model calculations are required to solve the present uncertainties in the overall interpretation of the data.  This underlines the importance of a new attempt to determine the geometry from radio polarization data. Then detailed model calculations starting from this geometry can attempt to reproduce all the remarkable timing and spectral results obtained for this millisecond pulsar."
    },
    "0206/astro-ph0206048_arXiv.txt": {
        "abstract": "{ A new method is presented for recovering the abundances of Damped Ly $\\alpha$ systems (DLAs) taking into account the effects of dust depletion. For the first time, possible variations of  the chemical composition  of the dust are taken into account in estimating the depletions. No prior assumptions  on the  extinction properties of the dust are required. The method requires a set of abundances measured in the gas and a set of parameters describing the chemical properties of the dust. A large subset of these parameters is determined from interstellar observations; the others are free   parameters for which an educated guess can be made. The  method is able to recover the  abundances of the SMC starting from SMC interstellar measurements apparently discrepant from the stellar ones. Application of the method to 22 DLAs with available [Fe/H] and [Si/Fe] measurements gives the following results: (1) the mean  metallicity of the corrected data is $< \\mathrm{[Fe/H]}> \\, \\simeq -1.0 $ dex, about 0.5 dex higher than that of the original data; (2) the slope of the [Fe/H] versus redshift relation is steeper for the corrected data ($m \\simeq -0.3$ dex) than for the original ones ($m \\simeq -0.2$ dex); (3) the corrected [Si/Fe] ratios are less enhanced, on  average, than those found in Galactic   stars of similar, low metallicity; (4) a   decrease of the  [Si/Fe] versus [Fe/H] ratios, expected by \"time delay\" models of chemical evolution, is found for the corrected data; (5) the [Si/Fe] ratios tend to increase with redshift once corrected; (6) consistency between [Si/Fe] and [S/Zn] measurements, two independent estimators of the $\\alpha$/Fe ratio, is found only for the corrected abundances. ",
        "introduction": "High resolution spectroscopy of   QSO  absorption line systems is a very powerful tool for probing structures in the early universe, including galaxies in the faint end of the luminosity function. With this technique, elemental abundances can be measured in     galaxies up to redshift $z \\approx 5$ and used to probe the early stages of chemical enrichment. In fact,    abundance measurements have been obtained for a relatively large number of damped Ly $\\alpha$ systems (DLAs), the class of QSO absorbers most clearly associated with intervening galaxies (Lu et al. 1996, 1998; Pettini et al. 1997, 1999, 2000; Prochaska \\& Wolfe 1999; Centuri\\'on et al. 1998, 2000; Molaro et al. 2000; Petitjean et al. 2000; Prochaska et al. 2001).  Abundance measurements can be rather accurate in DLAs: column densities of \\ion{H}{i} and of  metal species can easily be determined with errors $\\leq 0.1$ dex or even $\\ll 0.1$ dex; ionization corrections are generally negligible for most of the elements considered (Vladilo et al. 2001; see also Izotov et al. 2001). However, the final accuracy of DLA abundances can be severely affected by the effects of dust. Evidence of dust in DLAs comes from extinction   and  abundance studies. The reddening of QSOs lying behind DLA absorption is a clear sign  of dust extinction (Pei et al. 1991; Pei \\& Fall 1995). Evidence of elemental depletion   similar to that observed in the nearby ISM  (Savage \\& Sembach 1996) was first reported by Pettini et al. (1994) from the   behaviour of the Zn/Cr ratio in DLAs. The   enhancement of the Zn/Cr  ratio relative to the solar value was interpreted as being due to a different degree of incorporation into dust grains of zinc, a volatile element, and chromium, a refractory element. Similar results have been found   for other pairs of elements with differing degrees of affinity with dust, e.g. Zn/Fe. Recent abundance studies  give further support to the presence of  dust   depletion in DLAs (Hou et al. 2001, Prochaska \\& Wolfe 2002), indicating that   abundances taken at face value are of little use in understanding the chemical properties of DLA galaxies.   Various approaches have been used to circumvent  the problem of dust depletion in studies of DLAs abundances. One way is to  focus  on elements essentially undepleted, such as N (Lu et al. 1998, Centuri\\'on et al. 1998), O (Molaro et al. 2000, Dessauges-Zavadsky et al. 2001; Levshakov et al. 2002), S (Centuri\\'on et al. 2000), and Zn (Pettini et al. 1997, 1999; Prochaska \\& Wolfe 1999, 2001; Vladilo 2000). Another is to study DLAs with low dust content (Pettini et al. 2000; Molaro et al. 2000, L\\'opez et al. 2002). Although these studies do not depend, essentially, on the properties of the dust,  their application is limited to certain elements and to   particular sub-samples of DLAs. The only way to perform a general study of DLAs abundances is to quantify the effects of dust depletion.  By comparing the abundance ratios observed in DLAs with those measured in the interstellar gas of the Galaxy it is possible to derive some indication on the depletions of DLAs. Studies of this kind show a similarity to the depletions typical of Galactic  warm gas (Lauroesch et al. 1996; Kulkarni et al. 1997; Savaglio et al. 2000). In this type of work, two different types of intrinsic DLA abundances have been considered: solar abundances and abundances typical of metal-poor stars in the Milky Way. Clearly, it is rather difficult to disentangle  the intrinsic abundance patterns of DLAs with this type of approach.  A method for recovering the intrinsic  abundances in individual systems was presented by Vladilo (1998, hereafter Paper I). In that work the chemical composition and extinction properties of the dust were assumed to be the same  as those  of the dust in Galactic warm gas. The amount of dust extinction for a given level of metallicity, i.e. the dust-to-metals ratio, was  allowed to vary among DLAs. The dust-to-metals ratio of individual DLAs was  estimated assuming that zinc tracks the iron-peak elements and that the observed overabundances of Zn/Cr or Zn/Fe are  due entirely to differential depletion. A limitation of that method is that the dust composition may actually vary in different environments, even within the Milky Way (Savage \\& Sembach 1996). In addition, some recent investigations  indicate that zinc may not exactly follow other iron-peak elements (Primas et al. 2000; Umeda \\& Nomoto, 2001). According to Prochaska \\& Wolfe (2002), the assumptions adopted in Paper I may force the dust-corrected abundances to become closer to solar values. To overcome these limitations, a new procedure for estimating dust depletion in DLAs is presented in this paper. The equations of the procedure are based on a new expression for scaling dust  depletions    up or down according to changes of the physical/chemical conditions of the interstellar gas (Vladilo 2002, hereafter Paper II). This expression allows the dust chemical composition to vary according to changes of the dust-to-metals ratio and of the intrinsic    abundances of the medium. With this new formulation, deviations of the intrinsic Zn/Fe ratios from solar values can be consistently accounted for. At variance with Paper I, the present method does not make any assumption concerning the extinction properties of the dust in DLAs. The method is presented in Section 2, some examples of application  are described in Section 3, and the results are discussed and summarized in Section 4. ",
        "conclusions": "Chemical abundances of DLAs  probe the effects of nucleosynthesis in high \\ion{H}{i} density regions of the universe seen  at different cosmic epochs throughout a considerable fraction of the Hubble time. However, disentangling the effects of nucleosynthesis from those of dust depletion   is a rather difficult task. In this work a new method has been presented for quantifying dust   effects, with the aim of recovering the intrinsic abundances starting from the observed ones. There are several important differences with respect to the original method presented in Paper I and  previous work on the same subject.  The first difference concerns the definition of the dust-to-metals ratio. In Paper I this was defined in terms of extinction per unit level of metallicity, with the aim of estimating {\\em indirectly} the extinction of DLAs from the elemental depletions measured in each system. The broad agreement  of such indirect estimates with the direct,  but rather uncertain, estimates obtained from QSO reddening studies (Pei et al. 1991) is a positive test  of consistency of the results obtained in Paper I. However, in order to compare these two types of estimates it is necessary to assume that the extinction is proportional to the  amount of dust present in the system. This assumption is risky, since the extinction is  also determined by the geometry and size distribution of the dust grains. The present work  focuses only on the {\\em abundances} in the dust and in the medium (gas plus dust). In order to avoid any assumption on the {\\em extinction} of the dust, we have defined the dust-to-metals ratio as the fraction of atoms  in dust  form of an element Y chosen as a reference,  i.e. $r=N_\\mathrm{Y,dust}/N_\\mathrm{Y,medium}$.  The second difference is the way that elemental depletions are scaled up or down as a function of the dust-to-metals ratio. In the work by Kulkarni et al. (1997) and in Paper I the dust-to-metals ratio was allowed to change, but the abundances in the dust were kept constant. In this case, the fractions in dust of all elements are scaled up or down together. Therefore, all    fractions in dust are proportional to that of the reference element, i.e. $f_\\mathrm{X} \\propto r$. In the present work the fraction in dust of various elements have different dependence on $r$, according to the general law  $f_\\mathrm{X} \\propto r^{1+\\eta_\\mathrm{X}}$. This law  has been derived and calibrated in Paper II and is able to reproduce all the  types of depletion patterns observed in   the Milky Way interstellar gas (Savage \\& Sembach 1996) with a single set of $\\eta_\\mathrm{X}$ parameters.    The third difference is that, for the first time, elemental depletions are allowed to vary also as a function of the overall abundances of the medium. As shown in Paper II, the dependence of the fraction in dust on the intrinsic abundances must be of the type $f_\\mathrm{X} \\propto 10^{(\\varepsilon_\\mathrm{X}-1)  \\left[ { \\mathrm{X} \\over \\mathrm{Y} }\\right]}$. The introduction of the parameter   $\\varepsilon_\\mathrm{X}$ allows us to consider different behaviours of the dust composition. When $\\varepsilon_\\mathrm{X}=0$, the composition of the dust does not vary with the composition of the medium. This extreme case could be valid, for instance, if the number ratio of  two elements in the dust is {\\em uniquely} determined by the capability of these two elements to have chemical bonds. In general, however, we may expect that the abundance ratio X/Y in the dust will track the abundance ratio X/Y in the medium, in which case $\\varepsilon_\\mathrm{X}=1$. By considering both cases $\\varepsilon_\\mathrm{X}=0$ and $\\varepsilon_\\mathrm{X}=1$ we can take into account very different behaviours of the dust composition with respect to changes of the overall abundances, without entering into the details of the  complex physical processes  of dust formation, accretion and destruction.  As in Paper I, the present method can be applied to individual DLA systems. If the quality of the observational data is good enough, the method can be used for studying the abundances of distinct velocity components inside a given DLAs. The main steps of the procedure for recovering intrinsic abundances can be summarized as follows. First, the dust-to-metals ratio is  derived with  Eq. (\\ref{DustToMetalsRatio}). The most natural choice for the [X/Y] ratio to be used in that equation is the [Zn/Fe] ratio. In practice, the dust-to-metals ratio is derived  from the observed  [Zn/Fe] value and an educated guess of the intrinsic  [Zn/Fe] value. Since the equations of the new method  take into account the dependence of the abundances in the dust on the abundances in the medium, possible changes of the intrinsic [Zn/Fe] ratio in DLAs can now be  treated  self-consistently. Therefore the unknown value of the intrinsic [Zn/Fe] ratio  is now considered as an input parameter which, at variance with Paper I, may also assume values different from zero. At this point  the abundances of elements other than Zn can be recovered from the observed  ones by means of Eqs. (\\ref{X_H}) and (\\ref{DustCorrection}). The dust-corrected abundances can be obtained for different sets of input parameters.  In this way, we can test the stability of the results for different values of the $\\varepsilon_\\mathrm{X}$ parameter and of the intrinsic [Zn/Fe] ratio.  The    new method has been applied to correct the few existing   interstellar abundances measured with accuracy in a nearby galaxy, namely the SMC. The results of this exercise  support the validity of the  procedure. In fact, for all the abundance ratios considered (Fe/H, Mg/Fe, Si/Fe, S/Fe, Cr/Fe, Mn/Fe, and Ni/Fe) the  corrected interstellar abundances are in agreement with the intrinsic SMC abundances (known   from stellar data). The method yields consistent results for  the different lines of sight considered. The importance of the dust correction procedure is illustrated by the fact that the non-corrected interstellar data would  indicate an SMC metallicity $-2.0 \\leq \\mathrm{[Fe/H]} \\leq -1.2$, a strong enhancement of the [$\\alpha$/Fe]  ratios, and  a high degree of inhomogeneity of the SMC. All these indications are contradicted by stellar data.  The successful results obtained for a low-metallicity  galaxy such as the SMC are encouraging since the final goal of the method is   recovering the abundances of DLA systems, also characterized by a low metallicity level. There are other reasons why we feel justified in applying the dust correction method to  DLAs systems. In fact, the adopted scaling law (Eq. \\ref{ScalingLawExtragal}) can successfully model depletion patterns originated in a variety of interstellar environments, with different physical conditions (e.g. cold vs. warm), locations (e.g. Galactic halo vs. disk; SMC), and extinction properties (e.g., strong vs. weak 217.5 nm emission bump).   The results obtained by applying the present method to the sample of DLAs with available \\ion{H}{i}, \\ion{Si}{ii}, \\ion{Fe}{ii}, and \\ion{Zn}{ii} measurements can be summarized as follows.  The mean metallicity of the sample is $<\\mathrm{[Fe/H]}> \\simeq -1.0$ dex for the dust-corrected data, significantly  higher than that of the original data, $<\\mathrm{[Fe/H]}> \\simeq -1.5$ dex. The difference between observed and corrected metallicities can be as high as 1 dex in some cases. Also the [Fe/H] versus redshift relation in DLAs is affected by dust effects, the slope being $\\simeq -0.32$ and $\\simeq -0.17$ for the data with and without dust correction, respectively.  Also the Si/Fe ratio, the most frequently measured $\\alpha$/Fe ratio in DLAs, is significantly affected by depletion. Most of the dust-corrected [Si/Fe] ratios lie  below the median values representative of Galactic stars. The [Si/Fe] enhancement  at low metallicities  ([Fe/H] $\\simeq -1.5/-2.0$ dex) is typically [Si/Fe] $\\simeq +0.2$ dex. This value is somewhat lower than the [Si/Fe]$\\simeq +0.3$ dex plateau of the original [Si/Fe] data found by Prochaska \\& Wolfe (2002) at the same metallicities, suggesting that  some depletion  is   present also in that plateau.  Possible trends   of   [Si/Fe]   versus [Fe/H] and versus redshift are also affected by   depletion effects. The   corrected [Si/Fe] data  decrease with increasing [Fe/H] and   increase with $z_\\mathrm{abs}$. Both trends are qualitatively consistent with the expectations of the time delay between the injection of $\\alpha$-capture elements and that  of Fe-peak elements, predicted by models of chemical evolution. The uncorrected [Si/Fe] ratios  do not show such  trends.  An important  test  of the dust correction method is provided by  the comparison between   Si/Fe   and S/Zn data. In fact, the S/Zn ratio is an indicator of the  $\\alpha$/Fe ratio  which does not require dust correction to be applied since both S and Zn are essentially undepleted in the interstellar gas. The   [Si/Fe] ratios  show a   behaviour similar to that of the [S/Zn] ratios only if the corrections are applied. This result supports the validity of the method.  A detailed interpretation of the intrinsic DLA abundances is quite complex since DLAs   represent a heterogeneous, and probably biased, collection of interstellar regions associated with different types of galaxies. A detailed comparison of the   dust-corrected abundances with the  predictions of chemical evolution models of different types of galaxies will be presented in a separate paper (Calura et al. 2002)."
    },
    "0206/astro-ph0206424_arXiv.txt": {
        "abstract": "We present imaging and spectroscopic observations of the gravitationally lensed arcs in the field of RX~J1347.5$-$1145, the most X-ray luminous galaxy cluster known. Based on the detection of the \\oii\\ $\\lambda$3727 emission line, we confirm that the redshift of one of the arcs is $z = 0.806$. Its color and \\oii\\ line strength are consistent with those of distant, actively star forming galaxies.  In a second arc, we tentatively identify a pair of absorption lines superposed on a red continuum; the lines are consistent with \\ion{Ca}{2} $\\lambda$3933 (K) and \\ion{Ca}{2} $\\lambda$3968 (H) at $z = 0.785$.  We detected a faint blue continuum in two additional arcs, but no spectral line features could be measured.  We establish lower limits to their redshifts based on the absence of \\oii\\ emission, which we argue should be present and detectable in these objects.  Redshifts are also given for a number of galaxies in the field of the cluster. ",
        "introduction": "%  Gravitational lensing by galaxy clusters serves as a powerful probe of cosmological structure.  The lensing phenomenon provides information on both the mass distribution of the lensing cluster and the nature of the background population of faint field galaxies (e.g., Smail et al. 1993; Fort \\& Mellier 1994). Measurements of cluster mass also place useful constraints on the nature of dark matter and the cosmological parameter $\\Omega$ (e.g., Mellier, Fort, \\& Kneib 1993; Fort \\& Mellier 1994; Crone, Evrard, \\& Richstone 1994, 1996; White \\& Fabian 1995). The arcs that are often seen in deep images of galaxy clusters are the sheared and magnified images of faint, background galaxies (Paczy\\'nski 1987). The amplification provided by gravitational lensing enables spectroscopic studies of these faint galaxies, which would otherwise be extremely difficult, if not impossible. The spectra of the arcs yield information on the redshifts and stellar populations of the lensed galaxies (e.g., Smail et al. 1993; B\\'{e}zecourt \\& Soucail 1997; Ebbels et al. 1998; Hall et al. 2000; Campusano et al. 2001). The redshifts of the arcs constrain models of the the cluster potential and provide robust estimates of the total cluster mass.  RX~J1347.5$-$1145, at a redshift of 0.451, is the most luminous X-ray cluster known, with an X-ray luminosity in excess of 10$^{45}$ \\lum\\ (Schindler et al. 1995, 1997; Ettori, Allen, \\& Fabian 2001). The mass estimates based on its X-ray properties (Schindler et al. 1995), the Sunyaev-Zel'dovich effect (Pointecouteau et al. 1999), weak-lensing models (Fischer \\& Tyson 1997), and strong-lensing models (Sahu et al. 1998; Cohen \\& Kneib 2002) have yielded discrepant results for the total mass of this cluster. Cohen \\& Kneib (2002) speculate that we may be witnessing the merging of two clusters along a direction perpendicular to our line of sight.  Strong constraints can be placed on the cluster mass based on the redshifts of lensed background galaxies. Schindler et al. (1995) discovered an arc system in RX~J1347.5$-$1145; it comprises of two bright arcs located $\\sim$35\\asec\\ from the central dominant galaxies, at diametrically opposite points along the North-South direction.  This was confirmed by the observations of Fischer \\& Tyson (1997).  For the bright northern arc (Arc~1), Sahu et al. (1998) reported the detection of an emission line plausibly identified with \\oii\\ $\\lambda$3727 at a redshift of 0.81.  The bright southern arc (Arc~4) showed a faint blue continuum but no spectral line features. The high-resolution {\\it Hubble Space Telescope}\\ images of Sahu et al. (1998) also revealed three additional arcs in this cluster (Arcs 2, 3, and 5). Figure~1 shows an $R$-band image of the central region of the cluster with identification for the five arcs.  This paper presents new photometry and spectroscopy of the arc system in RX~J1347.5$-$1145. We confirm the emission-line redshift previously reported for Arc~1, report the tentative detection of an absorption-line redshift for Arc~2, and give limits on the redshifts of Arcs 3 and 4.  These measurements are in good agreement with the predictions from the lensing models of this cluster.  We also give redshifts for a number of galaxies in the field of the cluster. ",
        "conclusions": "%   Our primary aim is to determine or place limits on the redshifts of the arcs in RX~J1347.5$-$1145.  The redshift for Arc~1 is relatively secure. Consistent with the study of Sahu et al. (1998), we detected a strong emission line at $\\sim$6728 \\AA\\ whose most likely identification is \\oii\\ $\\lambda$3727 at $z = 0.806$.  As discussed by Sahu et al. (1998), Ly$\\alpha$ $\\lambda$1216 or \\ion{C}{4} $\\lambda$1550 can be ruled out based on the photometric colors: the Lyman limit would render the $B_{J}-R$ color much redder than that reported by Fischer \\& Tyson (1997).  More directly, our spectrum, which extends to $\\sim$4500 \\AA, shows no sign of any continuum decrement.  Strong optical lines such as H\\bet, \\oiii\\ \\lamb\\lamb4959, 5007, or H\\al\\ can be ruled out trivially by the redshift (0.451) of the lensing cluster.  The detection of the \\oii\\ line allows us to deduce a few basic properties concerning Arc~1. The measured \\oii\\ EW of $\\sim$75--95 \\AA\\ compares well with values previously found in gravitationally lensed arcs seen toward other galaxy clusters (Ebbels et al. 1998).  It is somewhat larger than in typical nearby late-type galaxies (EW $\\approx$ 50 \\AA; Kennicutt 1992b), but lies in the upper end of the EW distribution seen in distant, faint galaxy samples (e.g., Colless et al. 1990; Hammer et al. 1997).  The observed $V-I$ color is also consistent with that expected for the faint blue galaxy population at $z \\approx 0.8$ (Forbes et al. 1996).  Although slit losses prevent us from measuring accurate total fluxes, we can use the observed \\oii\\ emission-line flux to set a lower limit to the formation rate of massive (ionizing) stars.  Kennicutt (1998) gives the following empirical relation between \\oii\\ luminosity ($L_{\\rm [O~II]}$) and star formation rate (SFR): \\begin{equation} {\\rm SFR} (M_{\\odot} \\, {\\rm yr^{-1}}) = (1.4 \\pm 0.4) \\times 10^{-41} \\, L_{\\rm [O~II]} \\,\\, ({\\rm erg \\, s^{-1}}). \\end{equation} \\vskip 0.3cm \\noindent This derivation assumes a Salpeter (1955) stellar initial mass function and solar metallicity.  For an observed $F_{\\rm [O~II]} > 5.5\\times10^{-16}$ \\flux\\ and cosmological parameters $H_0$ = 75 \\kms\\ Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_{\\lambda} = 0.7$, we obtain SFR $>$ 24 $M_{\\odot}$~yr$^{-1}$.  The above calculation accounts for a Galactic extinction of $A_B = 0.268$ mag (Schlegel, Finkbeiner, \\& Davis 1998) corrected using the extinction law of Cardelli, Clayton, \\& Mathis (1989).  To obtain the {\\it intrinsic}\\ SFR, we need to know the flux magnification factor due to the lensing.  According to the lensing model of Allen, Schmidt, \\&  Fabian (2002), the magnification factor of Arc~1 is 7.7 (R.~W. Schmidt, private communication). Thus, the true lower limit to the SFR for Arc~1 is $\\sim$3 $M_{\\odot}$~yr$^{-1}$.  This is comparable to the level of star formation activity in nearby gas-rich spiral galaxies (e.g., Kennicutt 1983) but is significantly lower than those obtained for the arcs in Abell~2218 (Ebbels et al. 1996) and Abell~2390 (B\\'{e}zecourt \\& Soucail 1997; L\\'emonon et al. 1998).  It is unclear whether slit losses alone can make up for the difference.  For Arc~2, we tentatively identify the pair of absorption lines at 7017 and 7084 \\AA\\ with \\ion{Ca}{2} $\\lambda$3933 (K) and \\ion{Ca}{2} $\\lambda$3968 (H) at a redshift of 0.785.  This interpretation is plausible considering (1) the redness of the continuum, which is suggestive of an old stellar population, and (2) the absence of any strong emission lines blueward of the absorption lines.  An old stellar population should show a more prominent 4000 \\AA\\ break than observed, but the shape of the spectrum redward of $\\sim$7200 \\AA\\ is too uncertain (due to telluric molecular absorption bands and residuals from subtraction of sky lines) to be definitive on this point.  For an elliptical galaxy at $z \\approx 0.8$, the observed $V-I$ color of 2.7 mag corresponds to a present-day restframe $V-I \\approx 1.2$ mag (Poggianti 1997), consistent with the colors of local elliptical galaxies (e.g., Fukugita, Shimasaku, \\& Ichikawa 1995).  Allen et al. (2002) recently used {\\it Chandra}\\ data to refine the mass model for RX~J1347.5$-$1145.  Adjusting their model to reproduce the redshift measurement of Arc~1 by Sahu et al. (1998), Allen et al. (2002) predict the redshifts of the other arcs.  For Arc~2, they give $z = 0.75 \\pm 0.05$, in excellent agreement with our value.  Allen et al. (2002) predict $z = 0.97 \\pm 0.05$ for Arcs~3 and 4.  Since both of these objects have featureless continua very similar in shape to that of Arc~1 --- indeed, they are {\\it bluer}, suggesting an even younger stellar population --- it is reasonable to expect that nebular emission should be present at a comparable, if not even greater, strength.  However, no significant emission feature is discernible at $\\sim$7340 \\AA, the expected location of \\oii\\ at $z = 0.97$.  Unfortunately, the quality of the spectrum in the region $\\sim$7300--7400 \\AA\\ is degraded by imperfect removal of the OH sky lines (see, e.g., Osterbrock \\& Martel 1992).  Nevertheless, a narrow emission feature with a strength comparable to that of the \\oii\\ line in Arc~1, and perhaps even a factor 2--3 weaker, would almost certainly have been detected in Arc~4, since the two objects have virtually identical continuum levels.  These arguments suggest that the redshift of Arc~4 is greater than 1.04, where we have taken the upper limit of our bandpass to be 7600 \\AA.  This redshift limit does not appear to be in serious conflict with the predicted value.  The faintness of Arc~3 makes its spectrum highly uncertain at the red end, and we limit our discussion to wavelengths \\lax 7200 \\AA, where the continuum is clearly detected and not severely affected by systematic effects.  The absence of any emission features with equivalent widths greater than $\\sim$10 \\AA\\ suggests that the redshift is likely to be larger than 0.93.  This is consistent with the value predicted by Allen et al. (2002)."
    },
    "0206/astro-ph0206338_arXiv.txt": {
        "abstract": "We use three-dimensional hydrodynamic numerical simulations to study phase transformations occurring in a clumpy interstellar gas exposed to time-dependent volumetric heating. To mimic conditions in the Galactic interstellar medium, we take a numerical model of a turbulent multiphase medium from Kritsuk \\& Norman (2002) computed in a periodic box with mean density $n_0=0.25$~cm$^{-3}$ and mean pressure $P_0/k\\sim10^{3.4}$~K~cm$^{-3}$. A second model with $n_0=1$~cm$^{-3}$ is also considered. Variations of the heating rate on a timescale of 1--10~Myr applied thereafter cause pressure variations in the gas and shifting of the thermal equilibrium curve in the phase plane. This stimulates mass transfer between the gas phases via thermal instability, converting 5--10~\\% of the thermal energy into kinetic energy of gas motions. The experiments demonstrate that recurrent substantial heating episodes can maintain turbulence at this level. Possible applications to the interstellar gas heated by variable far-ultraviolet background radiation produced by short-living massive stars are discussed. ",
        "introduction": "Thermal instability (TI), controlled by the interplay of volumetric heating and radiative cooling of optically thin gas, is one of the key physical processes operating in the interstellar medium (ISM) \\citep{field65,meerson96}. For a long time it was believed that an important manifestation of TI is the splitting of the ISM into multiple thermal phases that coexist in pressure equilibrium \\citep{pikel'ner68,field..69,mckee.77,mckee90,heiles01a}. Recent numerical studies extended this classical picture, showing that interstellar turbulence quite naturally appears as a by-product of the same phase transformations that provide the irregular, clumpy and filamentary density substructure in the ISM \\citep[hereafter KN02]{koyama.02,kritsuk.02}. In the absence of continuous driving, however, this turbulence decays as $t^{-\\eta}$, with $1\\lsim\\eta\\lsim2$ [KN02; cf. \\citet{maclow...98,stone..98}], leaving behind a ``fossil'' isobaric density distribution with no substantial velocity structure remaining at the scales of the density inhomogeneities [cf. \\citet{mccray..72}].  Since observations indicate the persistence of turbulence and substructure in the ISM on a variety of timescales and length scales, there must be some source of energy for turbulence support. A number of hydrodynamic and MHD mechanisms have been proposed to sustain turbulent motions within molecular clouds and other thermal phases of the ISM. These include winds from young stars \\citep{norman.80,franco.84}, photoionization-regulated star formation \\citep{mckee89}, supernova explosions \\citep{mckee.77,korpi...99,kim..01}, largescale external shocks \\citep{kornreich.00}, and differential rotation \\citep{richard.99,sellwood.99,wada.99}. They rely on a combination of energy deposition from star-forming activity, self-gravity, magnetic effects, or galactic rotation to feed the turbulence. These options do not exhaust all possible driving mechanisms, cf. \\citet{norman.96}. Motivated by the results of KN02 on TI-induced ISM turbulence, we study here the possibility of purely ``thermal'' support for interstellar turbulence. This mechanism, which is physically distinct from the local mechanisms listed above, can operate {\\em in addition} to them.  Most previous works focused specifically on TI implicitly assume the existence of stationary thermodynamic equilibrium. The ISM, however, is highly time-dependent \\citep{gerola..74,bania.80}. In particular, time variations can be caused by changes in the level of ISM heating. The main source of energy input for the neutral gas is background radiation in the far-ultraviolet (FUV) part of the spectrum \\citep{wolfire....95}. The FUV field is generally due to OB associations of quickly evolving massive stars that form in giant molecular clouds. For example, in the local ISM the expected FUV energy density undergoes substantial fluctuations on a wide range of timescales from $\\lsim10$ to $\\sim100$~Myr \\citep{parravano..02}. The focus of the numerical simulations presented in this Letter is on the hydrodynamic effects of the time-dependent heating of the ISM.   \\begin{figure*} \\epsscale{2.0} \\plotone{fi1.ps} \\vspace{0.3cm} \\caption{Snapshots of the gas density field (perspective volume rendering): ({\\em a}) Turbulent multiphase gas at $t=2$~Myr; ({\\em b}) relaxed state at a high heating rate with no cold phase present, $t=2.8$~Myr; ({\\em c}) violent relaxation to an equilibrium with reduced heating, (the seeds for the new population of cold clouds are forming along caustics in the stable warm phase, $t=3.24$~Myr); ({\\em d}) partially relaxed state at a low heating rate, $t=4$~Myr (20~pc box, $128^3$ grid points). The log density color coding is as follows: The most dense blobs, $n > 8$~cm$^{-3}$, are blue; the less dense gas, $n\\in[3, 8]$~cm$^{-3}$, is light blue to green; the warm gas at $n\\in[0.5, 3]$~cm$^{-3}$ is yellow to red, and the low-density gas ($n < 0.5$~cm$^{-3}$) is transparent. The figure is also available as an mpeg animation in the electronic edition of the {\\em Astrophysical Journal}. \\label{fig1}} \\end{figure*} ",
        "conclusions": "Our model is limited in many respects. First, it does not include effects of other driving mechanisms that can operate in parallel to thermal forcing. Second, by not including magnetic fields in the simulations, our description of ISM dynamics is incomplete. Third, the model does not have self-regulation, linking production of the cold phase to the heating rate via star formation efficiency \\citep{parravano88}. Feedback could probably damp the oscillations occurring in response to the time variations of the heating rate introduced ``by hand,'' but this should be demonstrated by an explicit calculation. Fourth, the grain photoelectric heating rate is not simply proportional to gas density. This would slightly modify the shape of the thermal equilibrium curve as it shifts from a low to a high state. Finally, the equation of state and the cooling function that we use do not take into account dynamics of ionization, recombination, formation of molecules, etc., which are decoupled from hydrodynamics in our treatment. This could overestimate the minimum temperature of the cold phase and thus underestimate its density and the rms Mach number.  Nonetheless, our main conclusion that time-dependent heating supports turbulence in thermally unstable multiphase ISM remains valid in spite of the above-mentioned model simplifications. Thermal forcing can potentially excite turbulent oscillations on a wide range of length scales because of the penetrative nature of FUV radiation. The level of thermally sustained turbulence as a function of length scale at a given point within the Galactic disk is controlled by the local mean gas density and by the power spectrum of incident FUV flux time variations. It will share the pattern of spatial inhomogeneity of FUV energy density within the disk.  The thermally unstable gas mass fractions in our low- and intermediate-density models ($f_G\\sim 10\\%-30$\\%) are in reasonable agreement with measurements by \\citet{heiles.02} as well as the fraction and morphology of the density distribution of the cold neutral phase (see \\citet{vazquez-semadeni...02} for a review). Both these results imply a substantial level of turbulence in the ISM. We suggest time-dependent heating as one possible mechanism for driving this turbulence."
    },
    "0206/astro-ph0206497_arXiv.txt": {
        "abstract": "We have conducted a survey for warm-hot gas, traced by \\ovi absorption in the spectra of 5 high-redshift quasars ($2.2<z<2.8$) observed with Keck I/HIRES. We identify 18 \\ovi systems, 12 of which comprise the principal sample for this work.  Of the remaining six systems, two are interpreted as ejecta from the QSO central engine, and four have ionization conditions affected by proximity to the background QSO. Of the 12 intergalactic \\ovi absorbers, 11 are associated with complex systems showing strong \\lya ($N_{\\mhi} > 10^{15.2}$ cm$^{-2}$), \\civ, and often other lower ionization species.  We do not detect any lines that resemble photoionized, enriched gas associated with the lowest density regions of the \\lya forest ($13.5<\\log N_{\\mhi}<14.5$). Not all of the systems lend themselves to a straightforward determination of ionization conditions, but in general we find that they most closely resemble hot, collisionally ionized gas found near regions of significant overdensity.  The extent and gas density of the intergalactic \\ovi absorbing regions are constrained to be $L\\le 200$ kpc and $\\rho/\\bar{\\rho}\\ge 2.5$.  This was calculated by comparing the maximum observed \\ovi linewidth with the broadening expected for clouds of different sizes due to the Hubble flow.  For the median observed value of the Doppler parameter $b_{\\movi}=16$ km/s, the inferred cloud sizes and densities are $L\\sim 60$ kpc and $\\rho/\\bar{\\rho}\\sim 10-30$.  The clouds have at least two distinct gas phases.  One gives rise to absorption in photoionized \\civ and \\siiv, and has temperatures in the range $T=20,000-40,000$ K, and overdensities of $\\rho / \\bar{\\rho}\\ge 100$. The second phase is traced only in \\ovi absorption.  Its temperature is difficult to constrain because of uncertainties in the nonthermal contribution to line broadening.  However, the distribution of upper limits on the \\ovi, \\civ, and \\siiv temperatures indicates that the \\ovi thermal structure differs from that of the other ions, and favors higher temperatures where collisional ionization would be significant.  The \\ovi systems are strongly clustered on velocity scales of $\\Delta v=100-300$ km/s, and show weaker clustering out to $\\Delta v = 750$ km/s. The power law slope of the two-point correlation function is similar to that seen from local galaxy and cluster surveys, with a comoving correlation length of $\\sim 11h_{65}^{-1}$ Mpc. The average Oxygen abundance of the \\ovi systems is constrained to be $[O/H]\\ge -1.5$ at $z\\sim 2.5$, about 10 times higher than the level observed in the general IGM.  Two production mechanisms for the hot gas are considered: shock heating of pre-enriched gas falling onto existing structure, and expulsion of material by supernova-driven galactic winds.  Comparison between the observed numbers of \\ovi systems and expectations from simulations indicates that infall models tend to overproduce \\ovi lines by a factor of $\\sim 10$, though this discrepancy might be resolved in larger, higher-resolution calculations.   Known galaxy populations such as the Lyman break objects are capable of producing the amount of \\ovi absorption seen in the survey, provided they drive winds to distances of $R\\sim 50$ kpc. ",
        "introduction": "Within the last decade, a picture of the evolving intergalactic medium has emerged whereby the growth of baryonic structure is described through the collapse of gravitational instabilities \\citep{cen1994,miraldaescude1996,hernquist1996,zhang1995,petitjean1995}. According to this model, baryonic gas exists in several different states.  At high redshift, most of the gas is found in the \\lya forest, which is generally distributed and relatively cool at $T\\sim 10^4$ K, its temperature governed by photoionization heating.  Beginning at $z\\sim 2.5-3$, an increasing fraction of the baryons undergo a period of shock heating as they fall onto large-scale structure.  The cooling timescale for this shock-heated phase is long, so by $z=0$ as many as $30\\%$ of the baryons may accumulate in gas with temperatures between $10^5-10^7$ K \\citep{cen1999,dave2001,fang2001}. The remaining $70\\%$ of the baryons at the present epoch have either never been shock heated above $T\\sim 10^4$ K, or they have cooled much further into highly overdense structures near the junctures of filaments.  In these very dense environments the effects of local processes begin to play an important role.  This picture must be incomplete at some level, since a substantial fraction of the universe is already metal-enriched by $z=4$ \\citep{meyer1987,womble1996,cowie1995,tytler1995}, and the enrichment process is not included in models relying solely on gravitational instability. Models of metal absorption lines in a hierarchical scenario \\citep{haenhelt1996, hellsten1997, cen1999, dave2001,fang2001} have relied either on a very early (``Population III'') pre-enrichment phase or on relatively simple global recipes for calculating stellar feedback. Ongoing metal enrichment at the epoch where we observe the metal absorption systems may be important, and galactic winds (one of the possible enrichment mechanisms) have been observed at both low and high redshift \\citep{heckman2001b,franx1997,pettini2001}. Moreover there are hints that metal enriched gas at $z=3$ is turbulent at levels which require energy input only $10-100$ Myr prior to the epoch of observation \\citep{rauch2001}.  Two of the principal phases of the IGM have been extensively studied because they are easily observed in the absorption spectra of high redshift QSOs as the \\lya forest (caused by the cool filaments) \\citep{kim1997, kim2001,rauch1997,mcdonald2001} and Lyman limit/Damped \\lya systems (caused by the regions of highest overdensity) \\citep[]{prochaska2000}. However, the hot phase of the IGM with $T>10^5$ K is comparatively poorly understood, because at such high temperatures the collisional ionization of Hydrogen becomes significant, rendering \\lya less effective for tracing structure.  A budget of the content of the IGM based on Hydrogen absorption alone will therefore underrepresent the contribution of hot gas to the baryon total.  A more accurate account of the hot phase may be made using species with higher ionization potential than Hydrogen.  The \\ovi 1032/1037\\ang doublet has long been recognized as a prime candidate for this purpose for several reasons \\citep[]{chaffee1986,dave1998,rauch1997}. First, the intergalactic abundance of Oxygen is higher than that of any element other than Hydrogen and Helium.  Second, highly ionized Oxygen in the form of \\ovi, \\ovii, or \\oviii is among a small number of effective tracers for gas in the $T=10^5-10^7$ K range typical of shocked environments in cosmological simulations \\citep[]{cen1999,dave2001,fang2001}. Among these ionization states, only \\ovi is visible in ground-based optical spectra of QSOs, at redshifts above $z\\ge2$.  Further interest in \\ovi has revolved around its predicted effectiveness for tracing heavy elements in the very low density IGM - an environment very different from the shock-heated one described above.  At $T\\sim 10,000-40,000$ K, the gas in this diffuse phase is too cold for collisional ionization to produce highly ionized species such as \\ovi.  However, its density is sufficiently low (only a few times the mean) that \\ovi may be produced through photoionization from the intergalactic UV radiation field.  Based upon recent simulations, one expects observable levels of photoionized \\ovi to exist in \\lya forest lines with column densities in the range $13.5 < \\log N_{\\mhi} < 15.0$ \\citep[]{hellsten1998,dave1998}.  This \\ovi absorption can therefore probe the metal content of gas with densities below the range in which \\civ is most sensitive.  Statistical studies involving the pixel-by-pixel comparison of optical depths of \\lya and \\civ have provided some evidence of widespread enrichment of the IGM to even the lowest column densities \\citep{songaila1996, ellison2000}, and more recently very similar techniques have been used to infer the statistical presence of \\ovi associated with the forest \\citep{schaye2000}. But to date the number of direct metal line detections associatied with $\\log N_{\\mhi}<14.5$ \\lya lines is small, so the presence of warm photoionized \\ovi could help validate the assumption of widespread enrichment used in some of the simulations described above.  Considerable attention has attended the recent discovery of \\ovi absorbers in the local neighborhood ($z<0.3$) using HST/STIS and the FUSE satellite \\citep{tripp2000b,richter2001,sembach2001,savage2002}. Much of the early interpretation of these results has involved the difficult job of distinguishing whether particular absorption systems represent the warm photoionized, or hot collisionally ionized variety of \\ovi.  Early indications show that the low redshift population is mixed, with a slight majority of collisionally ionized systems. Regardless of the physical interpretations of these individual lines, it seems clear that the baryonic content of the Warm-Hot IGM may be significant at low redshift - possibly as much as 30\\% of $\\Omega_b$.  In this paper, we describe the results of a survey for \\ovi at high redshift, along the lines of sight to five bright quasars observed with the Keck I telescope and HIRES spectrograph.  Our survey covers the range $2.2<z<2.8$, which was chosen to optimize the tradeoff between signal-to-noise and contamination from the \\lya forest. We estimate the contribution of warm-hot gas to the total baryon budget at high redshift, and attempt to identify the physical environments of the \\ovi systems.  In Section 2 we describe our observing strategy and methods, and provide brief descriptions of the properties of individual \\ovi systems. In Section 3 we characterize the physical envronment of \\ovi absorbers, and calculate the contribution of warm-hot gas to the baryon total.  In Section 4 we discuss possible production mechanisms for the highly ionized gas in the context of cosmological simulations and galaxy feedback on the IGM.  Throughout, we assume a spatially flat universe with $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=65$ km/s/Mpc. ",
        "conclusions": "We have performed a survey for \\ovi absorption in the intergalactic medium at $2.2<z<2.8$ along the lines of sight to five quasars. Eighteen \\ovi systems were identified, 12 of which constitute the primary sample for this paper. The remaining six systems are located near their respective background quasars, and are either ejected from the central engine or affected by the strong local radiation field. At high redshift, blending from the \\lya forest is a significant limitation in \\ovi searches.  Accordingly, we have provided quantitative estimates of both the completeness of the sample and its contamination level.  The completeness - a measure of the fraction of true \\ovi systems that are detected (i.e. not blocked by the forest) - is found to be $f_c\\approx 41\\%$.  The contamination level, which measures the number of \\lya line pairs in the forest masquerading as \\ovi doublets, is estimated at $\\le 10\\%$ of our identified \\ovi systems.  The intergalactic systems we have detected share the following directly observable properties: \\begin{enumerate} \\item{They are located in highly overdense environments. These regions are characterized by strong \\lya absorption, either as a Lyman-limit system or as a collection of several clouds with $N_{\\mhi}\\sim 10^{15.5}$.  One \\ovi system is associated with a weak damped \\lya absorber.} \\item{\\civ absorption is seen in the environment of every \\ovi system, and other lower ionization species (\\siiv, \\cii, \\siiii, \\siii) are often present as well.} \\item{The detailed velocity structure of \\ovi does not generally match that of the other heavy elements.  The \\ovi profiles typically show less substructure than \\civ or \\siiv, and are sometimes offset from lower ionization lines by several hundred km/s. } \\end{enumerate}  Surprisingly, we have failed to detect any \\ovi lines associated with lower density regions of the \\lya forest.  Cosmological simulations \\citep{hellsten1998,rauch1997} predict that photoionized \\ovi should be an efficient tracer of metals in clouds with $13.5 < \\log(N_{\\mhi}) < 14.5$, yet we detect no intergalactic \\ovi lines near regions with \\hi column densities below $\\log(N_{\\mhi})=14.6$. The survey data were originally taken to the standards required to detect \\ovi in forest lines with [O/H]=-2.5, and we find several clean systems where \\ovi should be seen in the data at this level but is not.  However, there are many systems where we would not detect weak \\ovi due to forest blending or noise, and it remains to be seen whether this finding is statistically significant or is simply a reflection of cosmological scatter in the Oxygen abundance or ionizing radiation field. These questions will be addressed in forthcoming work.  To aid in our interpretation of the detected systems, we have performed a suite of ionization simulations using the CLOUDY software package.  Runs were made for the photoionization equilibrium case, and also at a range of fixed temperatures between $10^4-10^6$ K to investigate the onset of collisional \\ovi production.  From comparisons between the observations and ionization calculations, we have assembled a basic physical description of high redshift \\ovi systems that may be summarized as follows: \\begin{enumerate} \\item{Their physical extent and gas density may be conservatively constrained at $L\\le 200$ kpc and $\\rho/\\bar{\\rho}\\ge 2.5$.  This was calculated by comparing the maximum observed \\ovi linewidth with the broadening expected for clouds of different sizes due to the Hubble flow.  For the median observed value of $b_{\\movi}=16$ km/s, the inferred cloud sizes and densities are $L\\sim 60$ kpc and $\\rho/\\bar{\\rho}\\sim 10-30$.} \\item{They possess at least two distinct gas phases.  One of these gives rise to absorption in photoionized \\civ and \\siiv, and has temperatures in the range $T=20,000-40,000$ K, and overdensities of $\\rho / \\bar{\\rho}\\sim 500$.  This is the same variety of gas which gives rise to the common \\civ and \\siiv absorption seen in all high redshift quasar spectra.  The second phase is physically distinct, and traced only in \\ovi absorption.  Its temperature is difficult to constrain because of uncertainties in the nonthermal contribution to line broadening.  However, the distribution of $T_{\\rm max}$ shown in Figure \\ref{fig_b_temp} indicates that the \\ovi temperature structure differs from that of \\civ and \\siiv, and favors higher temperatures where collisional production of \\ovi would be significant.} \\item{They are strongly clustered on velocity scales of $\\Delta v=100-300$ km/s, and show weaker clustering signal out to $\\Delta v = 750$ km/s. The power law slope of the two-point correlation function is similar to that seen from local galaxy and cluster surveys, with a comoving correlation length of $\\sim 11h_{65}^{-1}$ Mpc, intermediate between galactic and cluster scales.  While the correlation at large velocities could be the signature of galactic winds, a simple geometric wind model does not accurately predict the shape of the TPCF over the entire velocity range where signal is seen.  We have argued that the signal in the \\ovi TPCF is dominated by the signature of large scale clustering.  This does not rule out the possibility that its shape is slightly modified by peculiar motions or outflows.} \\item{On average, they possess Oxygen abundances of $[O/H]\\ge -1.5$, about 10 times higher than the level observed in the general IGM.  This level of enrichment probably requires additional metal input above the level supposed to originate from Population III stars.  This conclusion assumes that the \\ovi absorbers contain $\\le 10\\%$ of the baryons at high redshift, i.e., that they are distinct from the more tenuous IGM that gives rise to the \\lya forest and contains $\\ge 90\\%$ of the baryons.} \\item{The integrated mass of warm-hot gas at $z\\sim2.5$ amounts to $\\ge 0.5\\%$ of $\\Omega_b$.  This may be compared with the latest estimates of \\citet{savage2002}, who measure $\\Omega_{\\rm WH}$ at low redshift to be only $25\\%$ higher using identical assumptions. If the average metallicity of the universe increases with time, the difference is even less pronounced.  This relatively weak evolution may indicate that \\ovi is seen only during a short phase in the cooling history of pre-heated gas.  In this case, the balance of the shock heated IGM, integrated over time, could reside in gas with $T>10^6$ K, which has long cooling timescales and is not efficiently traced by \\ovi.} \\end{enumerate}  Based on the absorption data alone, we cannot make a strong distinction between the scenario where \\ovi is produced in wind-induced shocks, and the scenario where it is produced in accretion shocks due to structure formation.  The high average metallicity we have measured in the \\ovi absorbers seems to point towards the wind hypothesis, though we cannot measure metallicities on a system-by-system basis. The overdensities we associate with the observed \\ovi systems are between those observed in the \\lya forest and collapsed structures.  At the low end of the allowed density range, \\ovi with properties similar to those observed can be produced through either collisional ionization or photoionization, though some systems can only be explained by collisional processes even at low density.  At higher densities than $n_H\\sim 3\\times 10^{-4}$ collisional ionization is required to explain the strength of \\ovi relative to other ions for solar relative abundances.  We have compared the \\ovi absorbers to structures of similar overdensity and temperature in cosmological simulations to test the plausibility of the accretion hypothesis.  Using estimates of the number density and cross section of these structures, we find that the number of \\ovi detections predicted for the survey is too high by a factor of $\\sim 15$.  However, the comparison is not always straightforward because simulations in the literature typically contain either sufficient spatial dynamic range to resolve the scales of interest, or a large enough volume to minimize cosmic variance - but not both.  The generation of simulations currently running should be able to address this question more accurately, to distinguish whether this discrepancy is physical or numerical.  A critical assessment of the wind model is also difficult, as current cosmological simulations are not capable of treating such a complex processes in full physical detail. However, we have compared our results with the most well studied population of wind-producing galaxies at high redshift, the Lyman break galaxies.  Using current estimates of the LBG comoving number density, we find that they are capable of producing all of the observed \\ovi absorption if each LBG drives winds to a radius of $\\sim 41$ kpc - similar to the size inferred from our pathlength analysis. Furthermore, the time required to drive winds to this distance at $v_{\\rm wind}=100-200$ km/s is in agreement with recent estimates of the star formation ages of the LBGs.  While this coincidence does not constitute a direct connection between LBGs and \\ovi absorbers, it does demonstrate that known galaxy populations could plausibly give rise to the amount of \\ovi seen in our survey.  Ultimately, much of the warm-hot gas that has been suggested as a baryon reservoir at low redshift may be hotter than $10^6$ K, and hence be undetectable in \\ovi. The fact that $T_{\\rm max}$ for \\ovi spans the entire range from $10^5 < T < 10^6$ K without any clear peak suggests that the high temperature limit for \\ovi lines results from further ionization of the Oxygen, rather than the actual detection of a maximum gas temperature.  Gas which is shocked to temperatures of $2\\times 10^6$ K (near the peak of the predicted warm-hot temperature distribution from Dav\\'e et al 2001) at $\\rho/\\bar{\\rho}\\sim 5$ would have a cooling timescale of $\\tau_{\\rm cool}\\sim 1.5$ Gyr, which amounts to $60\\%$ of the Hubble time at $z=2.5$.  At the same density, gas in the \\ovi temperature range cools $\\sim 30$ times faster, owing to its lower initial energy, and its location at the peak of the cooling curve. If this is the case, then \\ovi absorption would trace the ``tip of the iceberg'' with respect to the total amount of warm-hot gas produced over cosmic time.  In principle, the prospect of detecting the rest of the gas is better in \\ovii or \\oviii X-ray absorption, but the number of extragalactic objects bright enough for high resolution X-ray absorption spectroscopy on current instruments is small, and all are at low redshift.  The indications from the \\ovi data, along with a handful of \\ovii and \\oviii detections, suggest that this hot gas exists, but it may be some time before it can be observationally characterized in a statistically robust sense.  Since the lifetime of gas in the \\ovi state is short, sources of energy input are needed to produce and maintain the high state of ionization in these systems.  Hence, the primary utility of \\ovi may not be for measuring the total content of the warm-hot intergalactic medium, or tracing the metal content of the lowest density regions of the forest, but rather for probing physics at the interface between galaxies and the IGM.  This connection has been suggested by the association of galaxy groups and \\ovi absorption at low redshift \\citep{savage2002}; an analogous connection at high redshift could aid in the characterization of processes thought to have significant impact on the thermal and chemical history of the IGM."
    },
    "0206/astro-ph0206174_arXiv.txt": {
        "abstract": "Andersson et al.\\ and Bildsten proposed that the spin of accreting neutron stars is limited by removal of angular momentum by gravitational radiation which increases dramatically with the spin frequency of the star.  Both Bildsten and Andersson et al. argued that the $r-$modes of the neutron star for sufficiently quickly rotating and hot neutron stars will grow due to the emission of gravitational radiation, thereby accounting for a time varying quadrupole component to the neutron star's mass distribution.  However, Levin later argued that the equilibrium between spin-up due to accretion and spin-down due to gravitational radiation is unstable, because the growth rate of the $r-$modes and consequently the rate of gravitational wave emission is an increasing function of the core temperature of the star.  The system executes a limit cycle, spinning up for several million years and spinning down in less than a year.  However, the duration of the spin-down portion of the limit cycle depends sensitively on the amplitude at which the nonlinear coupling between different $r-$modes becomes important.  As the duration of the spin-down portion increases the fraction of accreting neutron stars which may be emitting gravitational radiation increases while the peak flux in gravitational radiation decreases.  Depending on the distribution of quickly rotating neutron stars in the Galaxy and beyond, the number of gravitational emitters detectable with LIGO may be large. ",
        "introduction": "Accretion onto the surface of a neutron star can in principle spin up the rotation of the neutron star until the spin frequency equals the Kepler frequency of the inner edge of the disk.  In low-mass x-ray binaries, the disk is thought to extend to stellar surface so the maximal frequency that the neutron star can achieve exceeds 1 kHz. However, the observed and inferred spin frequencies of neutron stars in low-mass x-ray binaries (LMXBs) are clustered around 250--500 Hz \\citep[e.g][]{1998ApJ...501L..89B}; the millisecond X-ray pulsars SAX~J1808.4 and XTE~J1751 have slightly higher frequencies of 402~Hz and 435~Hz \\citep{1998Natur.394..344W,2002IAUC.7867....1M}. Millisecond radio pulsars have been discovered with frequencies up to 640~Hz \\citep{1982Natur.300..615B}.  All of these limits are well below the Keplerian limit on the spin frequency of a neutron star, so an alternative explanation for the maximal observed spin frequency of neutron stars is required.  \\citet{1999ApJ...516..307A} and \\citet{1998ApJ...501L..89B} proposed that inertial modes (specifically the $r-$modes) inside the neutron may grow while generating gravitational waves (GW).  For sufficiently quickly rotating stars, GW carry away the angular momentum as quickly as it is deposited on the star by accretion.  \\citet{1999ApJ...517..328L} found that this proposed equilibrium between spin-up and spin-down is unstable.  A neutron star will execute a limit cycle \\citep[see][for additional discussion]{2000ApJ...534L..75A}.  It spins up for several million years and then quickly spins down emitting GW in less than a year. Only a small fraction of neutron stars is spinning down at any time; it is unlikely that any neutron star within the galaxy is currently spinning down, so none would be detected by LIGO.  However, the duration of the spin down depends sensitively on the assumed maximal amplitude of the $r-$mode.  \\citet{1999ApJ...517..328L} assume that the $r-$modes saturate when their amplitude is of order unity. \\citet{2001astro.ph.10487S}, \\citet{2001ApJ...549.1011W} and \\citet{Arra02} found that the saturation amplitude may be two to three orders of magnitude smaller; this may increase the duration of spin down to be greater than several thousand years.  This dramatically increases the number of neutron stars whose GW could be detected both by increasing the typical GW flux relative to the estimates of \\citet{1999ApJ...516..307A} and \\citet{1998ApJ...501L..89B} and the number of currently emitting sources relative to \\citet{1999ApJ...517..328L}.   This Letter explore the implications of saturation of $r-$modes in rapidly rotating neutron stars whose spins are accelerated by accretion.  \\S\\ref{sec:calc} will describe a series of straightforward calculations similar to those of \\citet{1999ApJ...517..328L} but with various values of the saturation amplitude (\\S\\ref{sec:spincalc}). Both \\citet{2000ApJ...534L..75A} and \\citet{2001IJMPhD.10..381A} have estimated the number of observable sources if the duty cycle corresponds to an $r-$mode saturation amplitude of unity.  The observed number of millisecond pulsars, the presumed descendents of LMXBS \\citep[e.g.][]{1991PhR...203....1B}, yields an estimate of the number of potential sources and their amplitudes (\\S\\S\\ref{sec:distcalc2}-\\ref{sec:distcalc3}) as a function of the duty cycle. \\S\\ref{sec:discuss} will outline some consequences of these results and the observational outlook. ",
        "conclusions": "\\label{sec:discuss}  If the spin of accreting neutron stars is indeed limited by the emission of gravitational radiation \\citep{1999ApJ...516..307A,1998ApJ...501L..89B}, low-mass X-ray binaries may be an important source for LIGO.  Although neutron stars may execute a duty cycle \\citep{1999ApJ...517..328L} of spin-up and spin-down which reduces the number of active sources at a given time, the sources that are active are typically much brighter than in a model where they emit constantly.  Unfortunately, the number of LMXBs that have been discovered actively accreting is too small ($\\sim 100$) to determine which effect dominates for duty cycles less than ten percent.  However, if low-mass X-ray binaries are assumed to be the exclusive progenitors of millisecond pulsars \\citep[see][ for an alternative]{1984JApA....5..209V}, the demographics of the millisecond pulsars in the Galaxy and beyond provides an estimate of the number of sources. Specifically, if the duration of the epochs when the neutron star is emitting gravitational radiation is greater than 10,000 years, several objects in the Galaxy will be above the detection thresholds of LIGO. For $\\tau_\\rmscr{on} \\sim$~10,000 year, these sources would be detectable throughout the Local Group.  Those objects in the Galaxy could also be detected from their X-ray emission which would be powered by gravitational radiation reaction; they would be GW-powered neutron stars.  For $\\tau_\\rmscr{on}$ much less than 10,000~years and much greater than one year, no sources are likely to be detectable even with an enhanced LIGO detector.  However, if the duration of gravitational wave emission per cycle is less than several years, several sources could be detected by an enhanced LIGO.  In this case, the sources will be located at a typical distance of 1 Gpc.  Depending on the nature of the $r-$mode instability in the cores of quickly spinning neutron stars, LMXBs may provide gravitational-wave beacons throughout the Galaxy and the Local Group or at cosmologically significant distances."
    },
    "0206/astro-ph0206032_arXiv.txt": {
        "abstract": "I generalize the inflationary flow equations of Hoffman and Turner to arbitrary order in slow roll. This makes it possible to study the predictions of slow roll inflation in the full observable parameter space of tensor/scalar ratio $r$, spectral index $n$, and running $d n / d \\ln k$.  It also becomes possible to identify exact fixed points in the parameter flow. I numerically evaluate the flow equations to fifth order in slow roll for a set of randomly chosen initial conditions and find that the models cluster strongly in the observable parameter space, indicating a ``generic'' set of predictions for slow roll inflation. I comment briefly on the the interesting proposed correspondence between flow in inflationary parameter space and renormalization group flow in a boundary conformal field theory. ",
        "introduction": "Inflationary cosmology\\cite{guth81,linde82,albrecht82} has become the dominant paradigm for describing the very early universe. Over the past twenty years, inflationary model building has been a prolific enterprise\\cite{lyth99}. Concurrently, cosmological observations have improved to the point that it is beginning to be possible to rule out models of inflation\\cite{kinney01,hannestad02}. Future observations, particularly the MAP\\cite{MAP} and Planck\\cite{Planck} Cosmic Microwave Background (CMB) satellites, promise to dramatically improve the situation in the near future\\cite{dodelson97,kinney98}. The key observational parameters for distinguishing among inflation models are the tensor/scalar ratio $r$, the scalar spectral index $n$, and the ``running'' of the spectral index, $d n / d \\ln k$, since different inflation models predict different values for these parameters.  It is desirable, however, to gain some insight into what the {\\em generic} predictions of inflation are without having to work within the context of some particular model. The standard lore of a small tensor/scalar ratio and nearly scale-invariant power spectrum is insufficient now that precision measurements of the CMB and large-scale structure are becoming a reality. Hoffman and Turner have proposed the method of inflationary ``flow'' to gain generic insight into the behavior of inflation models\\cite{hoffman00}. The flow equations relate the time derivatives of the slow roll parameters to other, higher order slow roll parameters. With a suitable choice of truncation, this makes it possible to study the dynamics of inflation models without having to specify a particular potential for the field driving inflation. In this paper we generalize the method from the lowest-order analysis of Hoffman and Turner and derive a simple set of flow equations which can be evaluated to arbitrarily high order, and which are in fact {\\em exact} in the limit of infinite order in slow roll. We perform a numerical integration of $10^{5}$ inflation models to fifth order in slow roll, and plot their predictions in the observable parameter space $(r,n,dn/d\\ln k)$. The predictions of the models cluster strongly in the observable parameter space, in fact even more strongly than was suggested by Hoffman and Turner. (However the qualitative character of their analysis is preserved at higher order in slow roll.) We emphasize that in this paper we limit ourselves to inflation driven by a single scalar field $\\phi$. The case of multiple-field inflation is in general much more complex.  This idea of flow in the inflationary parameter space has taken on different significance with recent ideas arising from the ``holographic'' correspondence between de Sitter space and boundary conformal field theories proposed by Strominger\\cite{strominger01a}. Particularly interesting are efforts to interpret flow in the space of slow roll parameters as renormalization group flow in a boundary conformal field theory\\cite{larsen02}. This raises the possibility that understanding the evolution of inflationary parameters is important not just for phenomenology, but for fundamental reasons as well.  The paper is organized as follows: Section \\ref{sechjreview} briefly reviews the very powerful Hamilton-Jacobi formalism for inflation. Section \\ref{secperturbations} discusses the generation of fluctuations in inflation and the relationship between the slow roll parameters and the observables in various exact and approximate solutions of the inflationary equations of motion. The hierarchy of flow equations is derived in Section \\ref{secsrhierarchy}.  Section \\ref{secfixedpoints} discusses the fixed points in the slow roll parameter space. Section \\ref{secevaluatingtheflowequations} discusses the details of the numerical solution. Section \\ref{secconclusions} presents conclusions. ",
        "conclusions": "\\label{secconclusions}  We have derived a set of inflationary ``flow'' equations based on the Hubble slow roll expansion of Liddle {\\it et al.}\\cite{liddle94} that is in principle exact when taken to all orders. These equations completely specify the dynamics of the inflationary system, so that any particular inflationary potential can be specified as a point in this parameter space. The past and future dynamics of the model are then determined by evaluating the flow of the parameters away from this point. It is possible to identify two classes of fixed points of the exact flow equations: power-law inflation, with $n = 1 - 2 r / (1 - r)$, and models with vanishing tensor/scalar ratio, $r = 0$. This latter class is unstable for $n < 1$ and stable for $n > 1$.  In practice, the flow equations must be truncated to some order and evaluated numerically, which was done to lowest order by Hoffman and Turner\\cite{hoffman00}. Extending the system of flow equations to higher order makes it possible to consider the running of the spectral index $d n / d \\ln k$ as well as $r$ and $n$. We perform a Monte Carlo integration of the flow equations to fifth order in slow roll, and show that the distribution of models in the parameter space of observables $r$, $n$ and $d n /d \\ln k$ is strongly clustered around particular values. Ninety percent of the models selected in the Monte Carlo converge to the observationally unacceptable asymptote $r = 0$, $n > 1.5$. The remaining models cluster around two classes of early-time ``attractor'', the first class at the $r = 0$ fixed point and the second with $r > 0$ and $n < 1$. Interestingly, the $r > 0$ attractor {\\em cannot} be identified with the power-law fixed point, since they generally have $d n / d \\ln k < 0$, and the variation in the spectral index vanishes at the fixed point. Evaluation of the models at very early times, $N \\gg 70$, indicates that the power-law fixed point is not an attractor at early times, since the models generically flow to the $r = 0$ line for large $N$. We therefore interpret the $r > 0$ ``attractor'' as simply an artifact of the fact that observable perturbations are generated relatively late in the inflationary evolution, when slow roll has begun to measurably break down. In addition, we see that power-law inflation is not in general an attractor for either early {\\em or} late times. At higher order, models much cluster more strongly than is suggested by the ``favored'' region of the parameter space derived by Hoffman and Turner. Also, models sparsely populate the regions labeled by Hoffman and Turner as ``excluded'' and ``poor power law'', suggesting that these categorizations do not generalize to higher order.  It is important to consider questions of generality with respect to both the choice of the order in slow roll $M$ and the choice of initial conditions for the Monte Carlo (\\ref{eqinitialconditions}). By ``closing'' the hierarchy of flow equations at finite order, we are implicitly limiting ourselves to a restricted class of potentials, although for $M = 5$, that class of potentials is large. However, models with potentials that contain features\\cite{starobinsky92,adams01} or for which the slow roll expansion is not convergent\\cite{wang97} will not be captured by solutions at finite order in slow roll. In addition, inflation might not be driven by only a single scalar field. The effect of different choices of initial conditions can be studied empirically, simply by trying different constraints on the space of initial conditions. Choosing ``looser'' initial conditions does not alter the characteristics of the result. Instead, models which fail to support sufficient inflation become much more numerous. Perhaps most importantly, absent a metric on the space of initial conditions, one should use caution when attempting to interpret these ``scatter plots'' statistically. We do not know how the initial conditions for the universe were selected! However, if observations determine that the relevant cosmological parameters lie outside the ``favored'' region, it will be an indication of highly unusual dynamics during the inflationary epoch.  Finally, we note an interesting recent body of literature connecting flow in inflationary models to a proposed ``holographic'' correspondence between quasi-de Sitter spaces and boundary conformal field theories (CFTs)\\cite{strominger01a,strominger01b,klemm01,bala01,bousso01,medved01,halyo02}. In particular, Larsen {et al.} have proposed a correspondence between slow roll parameters and couplings in the boundary CFT, interpreting flow in the inflationary parameter space as renormalization group flow in the associated CFT\\cite{larsen02,argurio02}. The fixed points at $r = 0$ are interpreted as ultraviolet ($n > 1$) and infrared ($n < 1$) fixed points in the renormalization group flow. In this picture, studying inflationary dynamics is equivalent to studying the structure of the underlying CFT. (It is not immediately clear, however, how one interprets the power-law fixed point in the context of the boundary CFT.)"
    },
    "0206/astro-ph0206204_arXiv.txt": {
        "abstract": "Temporal and spectral characteristics of prompt emission of gamma-ray burst (GRB) pulses are the primary observations for constraining the energizing and emission mechanisms. In spite of very complex temporal behavior of the GRBs, several patterns have been discovered in how some spectral characteristics change during the decaying phase of individual, well defined long ($>$ few seconds) pulses. In this paper we compare these observed signatures with those expected from a relativistically expanding, shock heated, and radiation emitting plasma shell. Within the internal shock model and assuming a short cooling time, we show that the angular dependence in arrival time from a spherical expanding shell can explain the general characteristics of some well defined long GRB pulses. This includes the pulse shape, with a fast rise and a slower decay, $\\propto (1+t/\\tau)^2$, where $\\tau$ is a time constant, and the spectral evolution, which can be described by the hardness-intensity correlation (HIC), with the intensity being proportional to the square of the hardness measured by the value of the peak, e.g. $\\Ep$ of the $\\nu F_\\nu$ spectrum. A variation of the relevant time scales involved (the angular spreading and the dynamic) can explain the broad, observed dispersion of the HIC index. Reasonable estimates of physical parameters lead to situations where the HIC relation deviates from a pure power law; features that are indeed present in the observations. Depending on the relative values of the rise and decay times of the intrinsic light curve, the spectral/temporal behavior, as seen by an observer, will produce the hard-to-soft evolution and the so called tracking pulses. In our model the observed spectrum is a superposition of many intrinsic spectra arriving from different parts of the fireball shell with varying spectral shifts. Therefore, it will be broader than the emitted spectrum and its spectral parameters could have complex relations with the intrinsic ones. Furthermore, we show that the softening of the low-energy power-law index, that has been observed in some pulses, can be explained by geometric effects and does not need to be an intrinsic behavior. ",
        "introduction": "The mechanism underlying the prompt $\\gamma$-radiation in gamma-ray bursts (GRBs) is still an unsolved puzzle.  There is, however, a growing consensus about some aspects of it.  The large energies and the short time scales involved require the $\\gamma$-rays to be produced in a highly relativistic outflow, an expanding fireball.  In the standard fireball model $\\gamma$-rays arise from shocks internal to the outflow at a distance of $R \\sim 10^{13}-10^{17}$ cm from the initial source. The episodic nature of the outflow causes inhomogeneities in the wind (or shells) to collide and thus creating the shocks.  These tap the bulk kinetic energy and transform it into random energy of leptons which radiate.  The dominant emission mechanisms are most probably non-thermal synchrotron \\citep{tavani, LP01} and/or inverse Compton emission \\citep{pan00}, but there have been other suggestions, for instance, thermal, saturated Comptonization \\citep{liang}.  The fundamental process of a burst is thus an individual shock episode which gives rise to a pulse in the $\\gamma$-ray light curve.  Superposition of many such pulses create the observed diversity and complexity of light curves \\citep{fish94}. The spectral and temporal characteristics of these pulses hold the key to the understanding of the prompt radiation of GRBs. However, there is no consensus on what effects lie behind the observed {\\it pulse} shapes and their temporal and spectral evolution.  The overall spectra of most GRBs can be described by a simple broken power law with a low and a high energy index, say $\\alpha$ and $\\beta$, and a break energy $E_b$. Often $\\alpha \\geq -2$ and $\\beta \\leq -2$, so that the $\\nu F_{\\nu}$ or $EF_E$ spectrum peaks at a photon energy $E_p\\sim E_b$.  The total light curves of GRBs, on the other hand, are very diverse and not readily describable by a simple formula. Nevertheless, many attempts have been made to decompose the complex light curves into pulses and analyze their characteristics \\citep{norris, lee1, lee2}. No simple patterns have emerged from these studies of the population as a whole.  However, some relations have emerged from investigations of GRBs with simple light curves; those described by a single pulse or a few, well separated pulses \\citep{kar95, RS00, RS02, BR01} (hereafter BR01). The pulse shapes and evolution of spectra seem to obey some simple relations. Motivated by these results, in this paper we explore possible explanations for these behaviors.  Several different possibilities exist.  The simplest scenario is to assume an impulsive heating of the leptons and a subsequent cooling and emission.  The rise phase of the pulse is attributed to the energizing of the shell which we will refer to as the {\\it dynamic time} and the decay phase reflects the cooling and its time scale. The instantaneous spectrum reflects the cooling of the lepton distribution.  The primary problem with this interpretation is that, in general, the cooling time for the relevant parameters is too short to explain the pulse durations and the resulting cooling spectra are in drastic disagreement with the above observed form \\citep{ghis}. A more plausible model is one where the pulse duration is set by the dynamic time of say the shell crossing, which could be much larger than the microscopic acceleration and/or emission-cooling times. In this case there is a continuous acceleration of particles during shell crossings; the acceleration and the cooling occur {\\it in situ} and simultaneously and give rise to the observed behavior.  The pulse shape then is a reflection of the energizing mechanism of the electrons.  A third possibility is that the above picture operates only during the rise phase of the pulse and that the decay shape is due to geometric and relativistic effects in an outflow with a Lorentz factor of $\\Gamma {\\gta} 100$.  The curvature of the fireball shell will make radiation, emitted off the line of sight (LOS, for short) delayed and affected by a varying relativistic Doppler boost,  due to the different light paths the photons have to travel.  The aim of this paper is to investigate to what extent, and how, the last model affects the observed light curve and spectral evolution during the individual pulses; in particular to determine whether the resultant behavior can explain the observed relations found for simple pulses mentioned above. We want to point out that the discussion of the observable signatures due to the curvature effect is independent of the process underlying the intrinsic pulses of radiation. Colliding shells and internal shocks are one example. However, other possibilities exist, for instance, as \\citet{lyu} pointed out, if the outflow is Poynting-flux dominated, the intrinsic radiation could be caused by a non-linear breakdown of large-amplitude electromagnetic waves at a distance of approximately $10^{14}$ cm from the progenitor. Furthermore, we emphasize that the description is for individual emission episodes, i.e. single pulses and one must bear in mind the possibility that these actually could consist of several heavily overlapping pulses; see further discussion in BR01 and \\citet{norris}. The observations relevant to our discussion will be described in \\S \\ref{sec:obsstat} and the appropriate time scales in \\S \\ref{sec:cond}. In \\S \\ref{sec:radiation} we derive the spectral and temporal structure expected in a simplified version of the proposed model. A more realistic model including both the dynamic and curvature effects is discussed in \\S \\ref{sec:broad}. Some other complications and caveats are discussed in \\S \\ref{sec:caveats} and a brief summary and discussion of the conclusions are given in \\S \\ref{sec:disc}.   In the following, primed quantities are evaluated in the comoving frame at rest with the outflowing material in the shock front. The rest frame will denote the inertial frame at rest with the progenitor.  The cosmological time dilation and spectral redshift, which are constant factors for individual pulses and bursts, will be ignored. ",
        "conclusions": "\\label{sec:disc}  We have examined the effect of differences in light travel time due to the curvature of the expanding shell and determined to what extent it can affect the width and shape of pulses and their spectral time evolution. The energy flux $F$ and the photon energies $E$ will be affected by the angle-dependent Lorentz-boost factor $\\cal{D(\\mu)}$; $F \\propto {\\cal{D}}^2$ and $E \\propto {\\cal{D}}$. The peak energy of the $\\nu F_\\nu$ spectra, $\\Ep$, will thus follow a hardness-intensity correlation (HIC) $F \\propto E^\\eta$ with $\\eta=2$. Furthermore, the decay phase of a pulse will follow the form $(1+t/\\ta)^2$. We  show that this effect should be important for a reasonable choice of parameters (Lorentz factor, burst energy, shell width etc.) and that these characteristics agree with the average behaviors found in pulses.  However, the curvature effect can not alone explain the large observed dispersion of the HIC index $\\eta$ . Cases with $\\eta$ largely different from $2$ we believe are produced by a finite dynamic time, $\\td \\neq 0$. The resulting spectral/temporal behavior depends mainly on the ratio $\\R = \\ta/\\td$. The intrinsic HIC (assumed to be a power law with an intrinsic index $\\gint$) will be revealed when $\\R < 0.5$ while the behavior expected from the curvature effect with a HIC index $\\eta=2 $ will dominate for $\\R > 2$. For intermediate $\\R$s the $ F-\\Ep$  relation will deviate  from a pure  power-law, having a more concave shape (in a $\\log - \\log$ plot). A general softening of the spectra with time, which has been observed, is also expected, independent of any changes in the intrinsic spectrum, and therefore independent of the physical environment where the pulses are produced.  An important conclusion of this work is that one must be very careful in the interpretation of the observed light curves and spectra, their parameters and evolution. This is because we have shown that the observed light curve will in most cases be different from the intrinsic one and the observed spectra will have a complex relation to the intrinsic ones. The spectra in the observer frame will be broader and, for instance, the low-energy power-law slopes $\\alpha$ will be softer than the intrinsic ones. In some cases, flat-topped spectra are produced which, in the observer frame, appear to have either $\\alpha$ or $\\beta= -1$. Furthermore, we also explain the occurrence of pulses whose $\\Ep$ track the flux up, and down, during the rise and decay phase, respectively, as well as the occurrence  of  pulses where the hardness declines monotonically independent of the rise and fall of the flux.  Ultimately, we wish to determine the characteristics of the {\\it intrinsic} emission, namely $F'(t)$, $\\Ep'(t)$, $\\gint$, and if possible the spectral power law slopes $\\alpha$ and $\\beta$. In addition, we want to determine the distance $R$ from the progenitor where the fireball emits the $\\gamma$-rays and to discern something about the shell width $\\Delta'$ and/or its spreading. Fits to the HIC and observed light curve will be able to reveal $\\R$ and $F'(t)$. Below we discuss the {\\it principle} diagnostics that can be made for three different situations. The value of the bulk Lorentz factor $\\Gamma$ remains an unknown parameter.  Case I. The observed HIC is a pure $\\eta=2$ power law ({\\it e.g.} pulses in Fig.\\ref{fig:2cases}): According to our model the curvature effect is dominant (It could, however, also be due to an intrinsic $\\eta=2$ HIC). The observed light curve will (asymptotically) follow equation (\\ref{ekv:flux}) with $d=2$, from which one can determine the value of $\\ta$. This time constant determines the distance at which the shell lights up: \\begin{equation} R = 2 \\ta \\Gamma^2v=6 \\times 10^{14} {\\rm cm} \\left(\\frac{\\ta}{1 \\rm {s}} \\right) \\left(\\frac{\\Gamma}{10^2} \\right)^2 \\beta \\end{equation} \\noindent The observed light curve can, in principle, be deconvolved,  with equation (\\ref{Ft_rel}) as the impulse response, to obtain $F'(t)$. A better knowledge of $F'(t)$, then gives a more accurate value of $\\R$ (and thereby $\\td$) from a fit to the HIC. Knowing $\\td$ one can put constraints on $\\Delta'/{\\beta_{\\rm sh}}'$ or ${\\beta_{\\rm sh}}' (\\Gamma_{\\rm s}/\\Gamma$). Furthermore, the observed, instantaneous spectra will be results of integrations of the intrinsic spectra along a $\\eta=\\gint$ power law. Combining this knowledge with the observations could reveal $\\Ep'(t)$ and $\\gint \\sim \\alpha_{\\rm obs}$ and possibly $\\alpha_{\\rm int}$.  Case II. The HIC is a pure power law with index $\\eta$, substantially different from $2.0$ (e.g. pulse in Fig. \\ref{fig:moreHICs}d): Here, $\\eta=\\gint$ and the light curve should reflect the intrinsic $F'(t)$ (smoothed somewhat by the curvature effect). The energy evolution follows $\\Ep'(t)$. Using $F'(t)$ a more thorough fit of the HIC can be made giving $\\R$, which gives an estimate of $R {\\beta_{\\rm sh}}'/ \\Delta'$ or ${\\beta_{\\rm sh}}' (\\Gamma_{\\rm s}/\\Gamma)$ (independent of $R$). The spectra arise from integrations along $\\eta=2$ which, depending on the details of the case, maybe provide a possibility to determine $\\alpha_{\\rm int}$.  Case III. Intermediate cases where S-curves are seen (e.g. pulses in Figs. \\ref{fig:moreHICs}a, b, and c): The low energy section of the HIC ({\\it i.e.} at late times) will follow $\\eta=2$ and gives the value of $\\ta$ and $R$. With this knowledge the light curve can be deconvolved and $F'(t)$ can be found. A fit to the HIC can now be made to find  $\\R$ [which gives $\\td$ and $\\Delta'$] and $\\gint$ which will be revealed from the early part, [which will allow the determination of $E'(t)$]. A corresponding softening of the spectra as described in the paper should be present.  BR01 found that in several GRBs, with two separable pulses, the HIC index varied less from pulse to pulse in a single burst as compared to its variation in different bursts. This requires that the pulses in multi-pulse bursts be produced in shocks created in a similar environment, with similar values of  $R$, $\\Gamma_{\\rm rel}$, $\\Gamma$, ${\\R}$, $\\gint$, $n$ and $B$. This could happen in a scenario in which the two long pulses are created as two similar shells catch up with a leading, slower, more bulky shell that has already been significantly decelerated due to interaction with the circumburst environment. Such pulses then occur approximately within the same environment, at roughly the same distance $R$ (therefore same $\\ta$) and $\\Gamma_{\\rm rel}$. This scenario also increases the value of $\\Gamma_{\\rm rel}$, which implies a higher magnetic field, radiative efficiency, and a minimum electron Lorentz factor, and a higher synchrotron peak frequency: \\begin{equation} h \\nu_{\\rm   s} = \\frac{3 e}{4 \\pi m_{\\rm   e} c} \\Gamma B \\gamma _{\\rm   e}^2= 5 {\\rm   eV} (\\Gamma_{\\rm rel}-1)^{2.5}\\left(\\frac{R}{10^{15}{\\rm   cm}} \\right)^{-1}, \\end{equation} where we have used the relations for B and $\\gamma_{\\rm e}$ described at the beginning of \\S 3. With $\\Gamma_{\\rm rel} \\sim 100$, $h \\nu_{\\rm s} = 500$ keV, so that the expected synchrotron spectrum will peak in the BATSE window and will require no additional boost, for instance, from Compton upscattering  as in the Synchrotron-Self-Compton model (SSC) \\citep{pan00}. This scenario is similar to  that of the external shock model normally proposed for the generation of the afterglows, which has difficulty to explain the prompt gamma-ray emission because of its high variability \\citep{Fen}. However, the GRBs discussed in this paper are smooth with few pulses and do  not exhibit the high variability of more complex bursts, so that this objection is not applicable."
    },
    "0206/astro-ph0206210_arXiv.txt": {
        "abstract": "We present the first mid-infrared (MIR) detection of a field brown dwarf (BD) and the first ground-based MIR measurements of a disk around a young BD candidate. We prove the absence of warm dust surrounding the field BD LP 944-20.  In the case of the young BD candidate Cha H$\\alpha$2, we find clear evidence for thermal dust emission from a disk. Surprisingly, the object does not exhibit any silicate feature as previously predicted.  We show that the flat spectrum can be explained by an optically thick flat dust disk. ",
        "introduction": "Brown Dwarfs (BDs) occupy the substellar mass domain. Having masses lower than 75 M$_{\\rm Jup}$,  they are unable to burn hydrogen steadily. Although their presence has been already predicted  in the sixties by \\cite{Kumar}, their low luminosity delayed their discovery until 1995, when \\cite{Nakajima}  announced the first detection of a BD orbiting the nearby M-dwarf star Gl229A. Recently, the large-scale near-infrared (NIR) surveys 2MASS and DENIS -- complemented by optical data -- substantially increased the number of known field BDs. Additionally, deep NIR surveys of star-forming regions  revealed hundreds of  young BD candidates.  In spite of the rapidly growing number of known BDs \\citep{Basri}, we do not know if they form like planets or like stars. Proposed scenarios  include the straightforward star-like formation via fragmentation and disk accretion  \\citep{Elmegreen}, the ejection of stellar embryos \\citep{Reipurth} from multiple systems and the formation in circumstellar disks like giant planets. The  presence of disks and their properties are crucial in distinguishing between the various scenarios: a truncated disk (size of a few AU) would support the ejected stellar embryo hypothesis, a non-truncated one is the sign of stellar-like accretion, while BDs formed like planets should have no dust around them.  In the case of BDs, NIR data are not necessarily a good tracer of disk emission because they are strongly affected by molecular bands of the cool BD atmosphere. Since the emission of warm (100 - 400 K) circumstellar dust peaks around 10 \\micron ,  mid-infrared (MIR) excess emission -- arising from dust grains close to the star -- is the best tool to search for  circumstellar disks. The MIR regime is best accessed by space-born telescopes, the last of which was the Infrared Space Observatory  (ISO), operating between 1995 - 1998. However, the majority of BDs has been discovered too late to be targeted by ISO.  Up to now, only few BDs with MIR excess are known. These objects, identified in the ISOCAM archive, are located in the Cha I or in the $\\rho$ Ophiuchi star-forming regions \\citep{Persi,Comeron98}. Their substellar nature has been deduced from comparing NIR and optical measurements to evolutionary models \\citep{Comeron98,Comeron}. \\cite{NattaTesti} proposed a model based on scaled-down disks around pre-main-sequence stars to explain the measured spectral energy distributions (SEDs). They followed the Chiang \\& Goldreich disk geometry which has been rather successfull in describing SEDs of T Tauri and Herbig Ae/Be stars \\citep{CG97,Natta, CG01}. The main assumption here is the flaring of the disk, which introduces a superheated surface layer, called the disk atmosphere (see Fig.~\\ref{Figure1}). \\begin{figure} \\plotone{f1.eps} \\caption{Cross sections of the flared and the flat disk model. The shaded area represents the optically thin superheated layer in the flared disk.  This region is the source of the silicate emission feature. The flat disk lacks this disk atmosphere. \\label{Figure1}} \\end{figure} This optically thin layer produces a strong silicate emission feature around 9.7 \\micron \\ (Si--O stretching mode). In an optically thick flat disk, no such feature is expected. Therefore, the presence or lack of such a feature is a strong constraint to any disk model.  In this paper we present results from our TIMMI2 MIR imaging campaign. Our aim was to detect MIR excess emission and thus to probe the presence of warm circumstellar dust around BDs. We targeted 8 very close field  BDs of various ages and a young BD candidate in the Cha I star-forming region. Our observations are the first data in the wavelength region of the silicate feature. ",
        "conclusions": "\\subsection{Field Brown Dwarfs} The non-detection of the 7 field BDs proves the lack of significant amount of warm dust around older field BDs. These data clearly show that the disk dissipation time is below a few 100 Myr, consistent with recent measurements of BDs in the young $\\sigma$ Orionis cluster \\citep{Oliveira}.  Even the detection of the closest target, the 475-650 Myr old LP 944-20, confirms this hypothesis. Compared to a simple blackbody with T$_\\star$=2300 K, R$_\\star$=0.1 R$_{\\odot}$, D=5 pc \\citep{Tinney}, it is clear that our measurements show no MIR excess, but the photospheric flux of the BD itself.   \\begin{figure} \\plotone{f2.eps} \\caption{ Modelled spectral energy distribution of a flat and a flared disk compared to the observations. The asterisks show the ISO measurements at 6.7 and 14.3 \\micron{} with 20\\% errorbars, while the squares are our TIMMI2 measurements at 9.8 and 11.9 \\micron{}. The estimated errors amount to 15\\%. The dot-dashed line indicates the contribution from the star, the dotted line is the flat disk emission. Their sum is the solid line fitting the observations. The dashed line shows the prediction of  the flared disk model.\\label{Figure2}} \\end{figure}   \\subsection{Disk Model for Cha~H$\\alpha$2} A black body with a temperature and radius corresponding to the classification of Cha H$\\alpha$2 fails to reproduce the observed fluxes that are an order of magnitude higher. This excess, previously observed by ISO \\citep{Persi}, was interpreted as emerging from a circumstellar disk \\citep{Comeron}. Our ground-based data corroborate the presence of a disk, and adds important constraints to the models.  The very low extinction towards Cha H$\\alpha$2 (A$_{\\rm I}<$0.4 mag, \\citet{Comeron}) argues for a nearly face-on disk, where a strong silicate emission feature would be expected from the superheated layer of a flared disk. Using the Chiang--Goldreich model \\cite{NattaTesti} predict a 10 \\micron \\ emission stronger than 40 mJy. In contrast, we measured fluxes of 17$\\pm$2 and 21$\\pm$3 mJy at 9.8 and 11.9 \\micron , respectively. The same flux densities at the peak and on the wing of the feature exclude the presence of any silicate feature.  In order to understand the reason of the model's deviation from the observations, we inspect its parameters and structure. The Chiang--Goldreich model consists of three major components: the star, the optically thin disk atmosphere and the optically thick disk interior (see Figure \\ref{Figure1}). In the MIR the stellar radiation is well approximated by a black body, while the optically thick disk emission is given by a power law $F_\\nu \\sim \\lambda^{5/3}$ \\citep{CG97}. A simple analytical formula is used to describe the optically thin disk atmosphere, which is producing the silicate emission \\citep{Natta}.  Changing the disk geometry (inner and outer radii, scale height, inclination)  is insufficient to explain the observed SED. Altering the optical properties (or composition) of the dust grains has a stronger effect: the absence of the silicate feature could be explained by the lack of this dust component or the presence of large grains (radii larger than 5 \\micron). However, we stress that the power law continuum $F_\\nu \\sim \\lambda^{5/3}$,  predicted by {\\em the flared model, does not fit our data}.  A much simpler and more straightforward solution is the assumption that the BD is surrounded by an optically thick flat disk. We assume a power law disk with a surface density of $\\Sigma \\propto R^{-3/2}$ and a temperature of $T \\propto R^{-3/4}$  which are typical of reprocessing and viscous disks \\citep{Shu}. Since this disk is entirely optically thick, its SED is independent of the dust properties. {\\em The model does not show any feature.  The continuum of a power-law flat disk has the observed slope}. In Fig. \\ref{Figure2} we compare the measurements with model predictions."
    },
    "0206/gr-qc0206061_arXiv.txt": {
        "abstract": "We study the evolution of the Weyl curvature invariant in all spatially homogeneous universe models containing a non-tilted $\\gamma $-law perfect fluid. We investigate all the Bianchi and Thurston type universe models and calculate the asymptotic evolution of Weyl curvature invariant for generic solutions to the Einstein field equations. The influence of compact topology on Bianchi types with hyperbolic space sections is also considered. Special emphasis is placed on the late-time behaviour where several interesting properties of the Weyl curvature invariant occur. The late-time behaviour is classified into five distinctive categories. It is found that for a large class of models, the generic late-time behaviour the Weyl curvature invariant is to dominate the Ricci invariant at late times. This behaviour occurs in universe models which have future attractors that are plane-wave spacetimes, for which all scalar curvature invariants vanish. The overall behaviour of the Weyl curvature invariant is discussed in relation to the proposal that some function of the Weyl tensor or its invariants should play the role of a gravitational 'entropy' for cosmological evolution. In particular, it is found that for all ever-expanding models the measure of gravitational entropy proposed by Gr{\\o}n and Hervik increases at late times. ",
        "introduction": "What is the entropy of a gravitational field? Following the insights of Hawking and Bekenstein \\cite{bekenstein,hawking}, an answer has been only been provided for a subset of situations characterised by static and stationary spacetimes with event horizons. In particular, black holes are found to be black bodies. They possess a well-defined entropy. They obey the laws of equilibrium thermodynamics and, when perturbed, their fluctuations obey the laws of close-to-equilibrium thermodynamics \\cite{sciama}. Yet, despite these developments it remains an unsolved problem whether an analogous gravitational entropy can be defined for non-stationary spacetimes of the sort that characterise most relativistic cosmologies that are considered as models for the past, present, and future of our visible universe. Unfortunately, a solution of this problem appears to involve the solution of a number of separate hard problems in gravitation physics. It requires a rigorous measure of the probability of different cosmological initial conditions to be found. It also requires us to be sure that we have accounted for all the possible contributions to the gravitational entropy: there may be quantum, classical, geometrical, topological, thermal, and dynamical contributions that need to be quantified and added \\cite{Hu}. The nature of the states of zero and maximum entropy also need to be defined and the influence of possible events like inflation must be fully accounted for entropically, and all must be suitably coordinate covariant. It has also been argued \\cite{davies74,davies83} that the whole concept of the arrow of time and thermodynamics relies upon the concept of a gravitational entropy, and hence leads all the way back to fundamental questions about the origin of time and the nature of the initial state of the universe.  Gravitation has a tendency to amplify small inhomogeneities in the distribution of matter. Yet, our present observations of the spectrum and temperature isotropy of the microwave background radiation show that the radiation was very close to thermal equilibrium when it was last--scattered, and that fluctuations from the Friedmann-Robertson-Walker (FRW) spacetime metric are only of order $10^{-5}$. If there does exist a gravitational entropy that grows with time to reflect the action of gravitational instability then it should have the property that it vanishes (or at least is small) for a homogeneous and isotropic metric and be maximal for a black hole of the same total mass. This type of observation has led some investigators to propose that any measure of cosmological gravitational entropy should be proportional to the deviation of the universe from the FRW model \\cite{bt}. Hence, the present gravitational entropy of the universe is in some sense small and so must have been even smaller in the past, perhaps providing some rationale for assuming the initial state of the universe to have zero entropy and hence be FRW up to minimal quantum gravitational fluctuations.  Based on this observation, Penrose \\cite{penrose} suggested that the Weyl curvature invariant, given by \\begin{equation} C^{\\alpha \\beta \\gamma \\delta }C_{\\alpha \\beta \\gamma \\delta }=R^{\\alpha \\beta \\gamma \\delta }R_{\\alpha \\beta \\gamma \\delta }-2R^{\\alpha \\beta }R_{\\alpha \\beta }+\\frac{1}{3}R^{2},  \\label{weyl} \\end{equation} should be very small near the initial singularity, and then grow thereafter behaving as a proper measure of the gravitational entropy. This in all its editions, is what we will call the \\emph{Weyl Curvature Conjecture} (WCC). Wainwright and Anderson \\cite{wa} had a different version of the WCC. They expressed the conjecture in terms of the relative magnitude of the Weyl and Ricci curvature invariants: \\begin{equation} P^{2}=\\frac{C^{\\alpha \\beta \\gamma \\delta }C_{\\alpha \\beta \\gamma \\delta }}{% R^{\\mu \\nu }R_{\\mu \\nu }}.  \\label{weyl1} \\end{equation} However, based on observations from two different universe models and taking into account that the entropy should scale with the volume, Gr\\o n and Hervik \\cite{wcc1}, \\cite{wcc2} suggested that a better measure of a gravitational entropy would be \\begin{equation} \\mathcal{S}=\\sqrt{|g|}P.  \\label{S} \\end{equation} Most other studies of the WCC \\cite{ra,rothman,gw,Bon85,Husain,pl,GCW} use the square of the Weyl scalar, $C^{\\alpha \\beta \\gamma \\delta }C_{\\alpha \\beta \\gamma \\delta },$ or $P,$ as the putative gravitational entropy for cosmological models. However, these candidates do not scale as the volume and so they do not capture correctly the dependence of a Weyl entropy on the volume.  We introduce the expansion-normalised shear: \\[ \\Sigma =\\frac{3}{2}\\frac{\\sigma _{\\mu \\nu }\\sigma ^{\\mu \\nu }}{\\theta ^{2}}. \\] The shear tensor is a symmetric and trace-free tensor defined by \\[ \\sigma _{\\mu \\nu }=\\theta _{\\mu \\nu }-\\frac{1}{3}\\theta (g_{\\mu \\nu }+u_{\\mu }u_{\\nu }). \\] Here, $u_{\\mu }$ is the normalised four-velocity of the comoving fluid; $% \\theta _{\\mu \\nu }\\equiv u_{\\mu ;\\nu }$ is the expansion tensor of the comoving fluid and $\\theta =\\theta _{~\\mu }^{\\mu }$ the volume expansion rate. The shear tensor tells us how anisotropic the universe is; in particular, $\\Sigma =0$ if and only if the universe is isotropic. From the generalised Friedmann equation ($8\\pi G=c=1$) \\beq \\frac 13\\theta^2=\\frac 12\\sigma^{\\mu\\nu}\\sigma_{\\mu\\nu}-\\frac 12{}^{(3)}R+\\rho+\\Lambda. \\label{friedmann}\\eeq Since the density, $\\rho$, is positive, $\\Sigma $ will be bounded by \\begin{equation} 0\\leq \\Sigma \\leq 1  \\label{sig} \\end{equation} whenever the spatial Ricci-scalar is non-positive: ${}^{(3)}R\\leq 0$. This inequality will be fulfilled in all of the homogeneous models,  except in the type IX and the Kantowski-Sachs closed universes.  One of the aims of this paper will be to investigate the quantity $% \\mathcal{S}$ for all homogeneous models, and determine whether or not it has the appropriate behaviour for a gravitational entropy. Hence, we will try to determine whether or not $\\mathcal{S}$ increases during the evolution of all homogeneous universes. We will also investigate the quantity $P$, because this quantity tells us something of the physical properties of the universe model under consideration and the general properties of the evolution of anisotropic universes. We will discuss, unless stated otherwise, only spatially homogeneous cosmologies with a non-tilted $\\gamma $-law perfect fluid. We also assume that the matter sources obey the strong energy condition; the effects from inflationary fluids are considered elsewhere \\cite{wcc2}. We will investigate all the possible spatially homogeneous models. Of special interest are the eight Thurston geometries \\cite{thurston97} which are related to the classification of topologies in three dimensions \\cite{thurston}. These models correspond to the Bianchi types in a rather interesting way \\cite {as,fik,Kodama1}. The possibility of compactification of these Thurston models has induced a renewed interest in the homogeneous models of Thurston and Bianchi type \\cite{BK1,BK2,Kodama2,sigBI}.  This paper is organised as follows. In section \\ref{FRW} we investigate all the models that have the FRW models as a special case. The remaining Thurston models are investigated in section \\ref{Thurston}, and in section \\ref{NonThurston} we investigate the remaining Bianchi models of non-Thurston type. We summarise our results in the final section.\\emph{\\ } ",
        "conclusions": "In this paper we have investigated the general behaviour of the Weyl curvature invariant for spatially homogeneous universes containing a perfect fluid as $t\\rightarrow 0$ and $t\\rightarrow \\infty $.  For all the models, except for the chaotic ones, the solutions asymptote in the past to points on the Kasner circle. The Kasner universe, which is a special vacuum case of the Bianchi type I model, has been treated exhaustively elsewhere \\cite {sigBI,wcc1}.  In the chaotic models, we saw that the Weyl curvature invariant $C^{\\alpha \\beta \\gamma \\delta }C_{\\alpha \\beta \\gamma \\delta }$, the invariant $P^2= \\frac{C^{\\alpha \\beta \\gamma \\delta }C_{\\alpha \\beta \\gamma \\delta }}{R^{\\mu\\nu}R_{\\mu\\nu}}$, and the invariant $% \\mathcal{S}=\\sqrt{|g|}P$ experience peaks and sinks. On the average, however, $\\mathcal{S% }$  increases steadily as the universe expands from the initial singularity.  The late-time behaviour was found to fall into one of five categories:  \\begin{enumerate} \\item  \\textbf{Ricci dominance.} This category includes models for which $P$ is decreasing at late times, indicating that the Ricci square dominates over the Weyl curvature invariant asymptotically.  \\item  \\textbf{Weyl-Ricci balanced.} In this category $P$ is approximately constant at late times. The Ricci and Weyl tensors are approximately proportional at late times.  \\item  \\textbf{Weyl dominance.} In this category $P\\rightarrow \\infty $ as $% t\\rightarrow \\infty $, but the expansion-normalized Weyl tensor, $\\mathcal{W}^2=C^{\\alpha \\beta \\gamma \\delta }C_{\\alpha \\beta \\gamma \\delta }/H^4$, will decrease. Except in the exceptional VI$_{-1/9}^*$ case, these models evolve towards vacuum plane-wave spacetimes, but the Weyl curvature invariant decreases more slowly than the Ricci square.  \\item  \\textbf{Extreme Weyl dominance.} In these models, $\\mathcal{W}$ $% \\rightarrow \\infty $ as $t\\rightarrow \\infty $.  \\item  \\textbf{Recollapsing models.} These models have no late-time behaviour, they recollapse after a finite time subject to the usual technical energy conditions required for recollapse \\cite{BGT}. \\end{enumerate}  \\begin{table}[tbp] \\centering \\begin{tabular}{|c|c|c|} \\hline Type & Matter & Category \\\\ \\hline\\hline I & $2/3<\\gamma \\leq 2$ & 1 \\\\ & $2/3<\\gamma \\leq 4/3 +\\pi_{\\mu\\nu}$ & 1 \\\\ & $4/3<\\gamma \\leq 2 + $ mag.field & 2 \\\\ \\hline II & $2/3<\\gamma<2 $ & 2 \\\\ \\hline III$_{LRS}$ & $2/3<\\gamma\\leq 3/2 $ & 2 \\\\ & $3/2 <\\gamma\\leq 2 $ & 3 \\\\ \\hline IV & $2/3 <\\gamma\\leq 2$ & 3 \\\\ \\hline V & $2/3 <\\gamma\\leq 2$ & 1 \\\\ \\hline VI$_0$ & $2/3 <\\gamma< 2$ & 2 \\\\ \\hline VI$_h$ & $2/3 <\\gamma<\\frac{2(1-h)}{1-3h} $ & 2 \\\\ & $\\frac{2(1-h)}{1-3h}\\leq\\gamma \\leq 2$ & 3 \\\\ \\hline VI$^*_{-1/9}$ & $2/3 <\\gamma \\leq 10/9$ & 2 \\\\ & $10/9<\\gamma \\leq 2$ & 3 \\\\ \\hline VII$_0$ & $2/3 <\\gamma < 1 $ & 1 \\\\ & $1<\\gamma\\leq 2$ & 4 \\\\ \\hline VII$_h$ & $2/3 <\\gamma\\leq 2 $ & 3 \\\\ \\hline VIII & $2/3 <\\gamma <4/5$ & 2 \\\\ & $4/5 <\\gamma \\leq 2 $ & 4 \\\\ \\hline IX & $2/3 <\\gamma \\leq 2 $ & 5 \\\\ \\hline KS & $2/3 <\\gamma \\leq 2 $ & 5 \\\\ \\hline I-IX & $0\\leq \\gamma <2/3$ & 1 \\\\ \\hline \\end{tabular} \\caption{Classification of the late-time behaviour. The matter content of the universe is labelled by the perfect fluid equation of state parameter, $\\gamma$; $\\pi_{\\mu\\nu}$ indicates the presence of trace-free anisotropic stresses and type I also includes the case of a pure magnetic field.} \\label{results} \\end{table} In the Table \\ref{results} we have summarised our results, and in the rightmost column we have indicated to which of the five categories each of the spatially homogeneous models belong. We note that category-3 cosmologies are not often discussed in the literature\\footnote{% This behaviour seems to have been noted for a special class of the VI$_{h}$ solutions in \\cite{GCW}.}. All of the plane-wave solutions of class B belong to this category. Note also that when we compactify the topology of the spatial sections \\cite{BK1,BK2}, most of the models in this category disappear in the compactification procedure; types IV and VI$_{h}$ cannot be compactified, type III must be locally rotationally symmetric (LRS), and VII$_{h}$ must be isotropic. However, some of the type III spacetimes (namely those LRS with $% \\gamma >3/2$) in category 3 can be compactified.  A general criterion to decide when a dynamical system will end up in category 3 can be given. If the late-time asymptote is a vacuum spacetime which is not conformally flat, but nevertheless has a zero Weyl curvature invariant, then the quantity $P$ will generically increase without bound. If, in addition, the state space of the corresponding dynamical system is also compact, then $\\mathcal{W}$ has to be bounded (and thus decreasing), and hence, the dynamical system will end up in category 3. The reason for this is as follows. For a conformally flat spacetime, both the Weyl curvature invariant and the Weyl tensor itself vanish. Hence, any linear perturbation of this conformally flat solution will produce linear corrections to the components of the Weyl tensor. Since the linear terms are the lowest-order terms in the Weyl tensor, quadratic terms will be the lowest-order terms in the Weyl curvature invariant. For non-conformally flat spacetimes, the Weyl curvature tensor will not be zero. Hence, the linear terms will not be the lowest-order terms in the perturbation. Thus, in the Weyl curvature invariant some of the linear terms will survive even though the lowest-order (constant) terms could vanish identically (which happens for the plane-wave solutions). The electric and the magnetic parts of the Weyl tensor are not zero for plane-wave solutions, they just happen to be of equal magnitude and cancel when the invariant is formed.  If the late-time asymptote is conformally flat, but not Ricci flat (hence, non-vacuum), then the system will typically end up in category 1. There is also a very small chance that the spacetime will end up in category 4 if the state space is non-compact. If the asymptote is both Ricci flat and conformally flat, all five categories are possible. If the state space is compact, then category 4 is ruled out. The fifth category depends to a certain extent on the geometry of the late-time attractor. For instance, inflationary spacetimes with $\\rho +3p<0$, are all governed by the ``no hair'' theorem which states that the late-time asymptote is a conformally flat non-vacuum spacetime. Hence, all ever-inflating cosmologies should end up in category 1.  In all of the models investigated, except for those in category 5, $\\mathcal{% S}$ increases at late times. Hence, $\\mathcal{S}$ behaves as expected by the WCC. It is not clear whether we should always expect a Weyl 'entropy' to increase if the initial conditions are close to some state of effective gravitational equilibrium. Of course, our study has only discussed spatially homogeneous universes and it remains to be seen whether small density and gravitational wave inhomogeneities evolve in a manner consistent with a thermodynamic interpretation of the Weyl curvature and whether the inclusion of significant inhomogeneities, black holes, and damping processes introduces significant new ingredients to the search for the elusive gravitational 'entropy'."
    },
    "0206/astro-ph0206026_arXiv.txt": {
        "abstract": "Motivated by the presence of numerous dark matter clumps in the Milky Way's halo as expected from the cold dark matter cosmological model, we conduct numerical simulations to examine the heating of the disk. We construct an initial galaxy model in equilibrium, with a stable thin disk. The disk interacts with dark matter clumps for about 5 Gyr. Three physical effects are examined : first the mass spectrum of the dark matter clumps, second the initial thickness of the galactic disk, and third the spatial distribution of the clumps. We find that the massive end of the mass spectrum determines the amount of disk heating. Thicker disks suffer less heating. There is a certain thickness at which the heating owing to the interaction with the clumps becomes saturates. We also find that the heating produced by the model which mimics the distribution found in Standard CDM cosmology is significant and too high to explain the observational constraints. On the other hand, our model that corresponds to the clump distribution in a $\\Lambda$CDM cosmology produces no significant heating. This result suggests that the $\\Lambda$CDM cosmology is preferable with respect to the Standard CDM cosmology in explaining the thickness of the Milky Way. ",
        "introduction": "Hierarchical clustering governed by cold dark matter (CDM) is widely believed as a cosmological scenario which is responsible for the growth of the structures in the universe. According to the hierarchical scenario, small dark matter halos should collapse earlier, but later fall into larger structures. The process of smaller halos being assembled into a larger halo does not always destroy the smaller ones, thus hierarchical structures are seen in many objects, such as clusters of galaxies.  Recent high-resolution simulations have successfully shown that hundreds of galaxy-size DM halos survive in clusters of galaxies \\citep{ghig1998, coli1999, klyp1999a}. A remarkable outcome of the high-resolution cosmological simulation in Standard CDM model by \\citet{moor1999} even shows that survival of substructures or satellites occurs not only on cluster scales, but also on galactic scales. They show that a galaxy of a similar size as the Milky Way should contain about 500 satellites, which is, however, much more than the number of the observationally identified satellites.  That many satellites should survive, was confirmed also by \\citet{klyp1999b} and \\citet{ghig2000}. Klypin et al. note that the results of the Standard CDM simulation are close to those of a $\\Lambda$CDM simulation with the same circular velocity function of substructures. This indicates that the prediction of the presence of many satellites is a general outcome of the hierarchical scenario and not particularly dependent on the cosmological models.  Compared with the observational results, these cosmological models yield a large number of the satellites, approximately a factor of~50 more than the number of satellites observed in the vicinity of the Milky Way.  \\citet{klyp1999b} suggest that the problem of the missing satellites could be resolved (i) by identification of some satellites with the high-velocity clouds observed in the Local Group \\citep{blit1999} or (ii) by considering dark satellites that failed to gather enough gas to form stars, because of expulsion of gas the supernova-driven winds or because of gas heating by the intergalactic ionizing background. The latter possibility implies that the halos of galaxies may contain substantial substructures in form of numerous invisible DM clumps.  A statistic of strong gravitational lensing is an approach to identify the dark clumps in the Milky Way. \\citet{chib2002} indeed finds evidence for the existence of the numerous satellites in the Milky Way.  If a great amount of the dark satellites exist within the Milky Way's halo, their interaction with the disk might cause disk heating.  In their Standard CDM model, \\citet{moor1999} found that a large fraction of the substructures have very eccentric orbits, so that they could cause resonant heating of the disk, and even heating by impulsive shocking owing to their penetration through the disk.  On the other hand, it is known that the Milky Way has a quite thin disk, whose scale height is about 200 pc. From the ``thinness'' \\citet{toth1992} have made an energetical analysis of the disk heating owing to accretion of matter, and derived the constrain that the Milky Way should have accumulated no more than 10 \\% of the disk mass within the past 5~Gyr. This estimation of the disk heating might, however, be too large because the actual interaction between the disk and satellites is more complicated. For example, a single satellite could dissolve before reaching the disk \\citep{huan1997} or the energy injected into the disk could excitate coherent warping motions rather than heat the disk \\citep{vela1999}. For the case of the interaction of the disk with many substructures in the halo, additional detailed numerical investigations are necessary. \\citet{font2001} have studied the case of the $\\Lambda$CDM cosmological model, and found that the substructures are not efficient perturbers for heating of the disk, since the masses of the clumps are lower than those of the clumps predicted in the Standard CDM model, and also because the clumps are located far away from the disk and seldom get near the disk.  The disk kinematics and dynamics is a good probe not only for examining the cosmological models, but much more generally for clarifying the halo substructure that is difficult to derive from direct observations. Therefore in this study we aim to investigate disk heating by dark matter (DM) clumps for a wide range of parameters.  Numerical experiments on the disk dynamics are not an easy task, especially when studying the vertical structure, because of the wide range of the dynamical scales among the different components in galaxies. The smallest scale is disk, with a scale height about 200~pc, while the largest scale is the dark halo extent of $\\gtrsim 100$ kpc. Many numerical studies consider disks of 700~pc ($0.2\\times$ the disk scale length) in thickness \\citep{vela1999, font2001}. The question is whether the heating rate obtained for such thick disks is the same as for thin disks like the real Galactic disk.  In this study we construct our initial galactic models following \\citet{kuij1995}, which is nearly in exact equilibrium and which allows us to set up disks as thin as the real Galactic disk. Several additional observational constraints are taking into account to build the galactic models. We study three physical effects that could affect disk heating; first the mass spectrum of the DM clumps, second the initial thickness of the galactic disks, and the third the spatial distribution of the clumps.  This paper is organized as follows. In section~2 we present the galaxy model which provides a very stable thin disk comparable to the real Milky Way disk. A model of a clumpy dark matter halo is presented in section~3. Numerical simulations of the interaction between disk and the clumps are specified in section~4. Section~5 presents the results of our numerical simulations on the disk heating by examining the effects on the mass spectrum of the clumps, the initial disk thickness, and the spatial distribution of the clumps. We summarize and discuss our results in section~6. ",
        "conclusions": "We have made a series of numerical simulations, examining heating of galactic disks that are embedded in a halo with many small DM clumps, as suggested by cosmological simulations. We have built up the initial galaxy model as precisely in equilibrium as possible, so that we could simulate a stable disk that is as thin as the Milky Way.  We have shown that with the mass spectrum $N(M)dM \\propto M^{-2}dM$, that is consistent with cosmological simulations, the massive end of the mass spectrum determines the amount of disk heating. The number of the clumps at the massive end could be only a few, so that the massive limit of the spectrum might fluctuate among galactic halos. This fact suggests a variety of disk thickness because of the random fluctuation in the number of the most massive clumps.  As a second result, the simulations demonstrate that thicker disks suffer less heating. It is possible that a disk with 700~pc thickness shows no significant heating whereas another disk with 200~pc thickness suffers significant heating in the same circumstances. This also means that there is a critical thickness at which the heating owing to the interaction with the clumps is saturated.  Finally, we have considered the relation to cosmology. Our model 2 mimics the clump distribution found in the Standard CDM cosmology \\citep{moor1999,klyp1999b}. The disk heating by the clumps has already been discussed by \\citet{moor1999}. They argued by a simple impulsive approximation that disk heating would be too high to explain the age-temperature relation for disk stars \\citep{wiel1974}. Our numerical simulations are supplementary to their analysis, and we confirm that the heating is significantly high. The thickness of 1~kpc at the solar radii is comparable to the present thick disk, but this heating takes place within 5~Gyr and even a young thin disk would have no chance to survive. On the other hand, our model 6 corresponds to the clump distribution of the $\\Lambda$CDM cosmology, similar to the one \\citet{font2001} have studied, but more closely representing the Milky Way. As found by \\citet{font2001}, this clump distribution causes no significant heating of the disk. These results suggest that the $\\Lambda$CDM cosmology is preferable with respect to the Standard CDM cosmology in explaining the thickness of the Milky Way."
    },
    "0206/astro-ph0206356_arXiv.txt": {
        "abstract": "s{Millisecond pulsar PSR J0218+4232 shows remarkable high-energy properties: very hard pulsed X-ray emission up to $\\sim 10$ keV and a likely detection of high-energy $\\gamma$-rays ($> 100$ MeV) with a soft spectrum. The relative phasing of the X-ray and $\\gamma$-ray profiles, however, was unknown. A recently performed Chandra (0.08-10 keV) observation of \\psr settled the phasing down to ~0.2 millisecond, and shows that the X-ray pulses are aligned with the $\\gamma$-ray pulses providing supporting evidence for our first detection of high-energy gamma-rays from this source. The preliminary results from a recent RXTE (2-250 keV) observation (27-12-2001 -- 7-01-2002) show significant pulsed emission, in a complex profile, up to $\\sim$ 20 keV. The composite high-energy spectrum of this millisecond pulsar is similar to the canonical spectrum of Unidentified Gamma-ray Sources (UGS), making fast ($<$ 3-4 ms) millisecond pulsars with low characteristic ages ($< 10^8$ years) good candidates for an UGS association and challenging targets for INTEGRAL observations. } ",
        "introduction": "\\psr is a 2.3 ms pulsar in a two day orbit around a low mass ($\\sim 0.2$ M$_{\\sun}$) white dwarf companion \\cite{nav}. Pulsed X-ray emission with a Crab-like double pulse profile has been reported from ROSAT 0.1-2.4 keV data \\cite{kuip1} and BeppoSAX MECS 1.6-10 keV data \\cite{min}. The pulsed spectrum as measured by the MECS appeared to be remarkably hard with a power-law photon index $0.61 \\pm 0.32$, harder than measured for any other radio pulsar. Furthermore, Kuiper et al. \\cite{kuip2} report the detection with EGRET of pulsed high-energy (0.1-1 GeV) $\\gamma$-ray emission from this millisecond pulsar. They also showed that the two $\\gamma$-ray pulses appeared to be aligned in absolute phase with two of the three radio pulses detected at 610 MHz. Unfortunately, the timing accuracies of the ROSAT and BeppoSAX observations were insufficient to construct X-ray profiles in absolute phase.  \\psr is also remarkable in that it is the only Crab-like ms pulsar with a large DC (unpulsed) fraction of $63 \\pm 13\\%$ in the ROSAT band below 2.4 keV \\cite{kuip1}, as well as a large DC fraction of $\\sim 50\\%$ in radio, systematically over the range 100-1400 MHz \\cite{nav}. The DC components as measured in the ROSAT and radio observations could be explained by emission from a compact nebula with diameter $\\sim 14\\arcsec$, but in both cases the indications were at the limit of the imaging capabilities. Assuming that the radio DC component is compact, combined with the measured very broad and structured radio pulse profile, Navarro et al. \\cite{nav} suggested that the magnetic field of \\psr is almost aligned with the rotation axis, the observer viewing the system under a small angle with respect to the rotation axis. Stairs et al. \\cite{stairs} determined the magnetic inclination angle analyzing radio polarization profiles. Their rotating vector model fits indicate that the magnetic inclination angle is indeed consistent with $0\\arcdeg$ ($8\\arcdeg\\pm11\\arcdeg$). Unfortunately, in their fits the line-of-sight inclination is unconstrained. If the DC component in X-rays is also compact, for the suggested geometry of a nearly aligned rotator and a small viewing angle, it can originate in the pulsar magnetosphere as well as from a heated polar cap of the neutron star.  The objectives of our Chandra and RXTE observations were: 1) To establish the spatial extent of the X-ray DC component, compact or extended (Chandra); 2) To construct an X-ray pulse profile which can be compared in absolute phase with radio profiles and the EGRET high-energy (0.1-1 GeV) $\\gamma$-ray profile (Chandra and RXTE); and 3) To investigate the characteristics of the pulsed emission for energies beyond 10 keV (RXTE).   \\begin{figure*}[t] \\hbox{\\hspace{-0.2cm}{\\psfig{file=idl_psrj0218+4232_hrci_map_lin_5x5_stgamma.ps,width=8.0cm,height=8.0cm}} \\hspace{0.2cm}{\\psfig{file=idl_psrj0218+4232_hrci_radprf_25_0.5.ps,width=8.0cm,height=7.5cm}}} \\caption{(left) Chandra 0.08-10 keV HRC-I image of a $10\\arcsec\\times 10\\arcsec$ region centered on the radio pulsar position of PSR J0218$+$4232. The radio position is marked with a `+' sign. The angular distance between the radio pulsar position and the X-ray centroid is $\\sim 0\\farcs 6$, consistent with the Chandra localization accuracy. (right) Radial distribution of HRC-I events using the optimum X-ray position as centre. Superposed as dotted line is the radial profile of the PSF. The dashed line indicates the background level derived from counts in the range 10 -- 25 arcsec from the centre. We have {\\bf no} indications for extended emission at $\\sim 1\\arcsec$ scales.\\hfill \\label{hrcispatial} } \\end{figure*} ",
        "conclusions": "From the recently performed deep Chandra and RXTE observations of \\psr we learned that:\\\\ $\\bullet$ \\hspace{0.2cm} There is no evidence for extended emission at X-ray energies at $\\sim 1\\arcsec$ scales.\\\\ $\\bullet$ \\hspace{0.2cm} The non-thermal X-ray pulses are aligned with 2 of the 3 radio pulses and with the two high-energy $\\gamma$-ray pulses strengthening the credibility of the first detection of pulsed high-energy $\\gamma$-ray emission from a (this) millisecond pulsar.\\\\ $\\bullet$ \\hspace{0.2cm} The total pulsed X-ray spectrum continues up to $\\sim 20$ keV into the soft $\\gamma$-ray domain.\\\\ \\noindent Future observations of this intriguing millisecond pulsar at soft $\\gamma$-rays by INTEGRAL and at high-energy $\\gamma$-rays by AGILE and GLAST are very important to determine its spectral characteristics in the $\\gamma$-ray regime in more detail and possibly to shed light on the origin of UGSs of which \\psr could be a canonical prototype."
    },
    "0206/astro-ph0206095_arXiv.txt": {
        "abstract": "We present first results from a 325~ks observation of the Seyfert 1 galaxy \\mcg\\ with \\xmm\\ and \\sax. The strong, broad, skewed iron line is clearly detected and is well characterised by a steep emissivity profile within $6r_{\\rm g}$ (i.e. $6GM/c^2$) and a flatter profile beyond. The inner radius of the emission appears to lie at about $2r_{\\rm g}$, consistent with results reported from both an earlier \\xmm\\ observation of \\mcg\\ by Wilms \\et and part of an \\asca\\ observation by Iwasawa \\et when the source was in a lower flux state. The radius and steep emissivity profile do depend however on an assumed incident power-law continuum and a lack of complex absorption above 2.5~keV. The blue wing of the line profile is indented, either by absorption at about 6.7~keV or by a hydrogenic iron emission line. The broad iron line flux does not follow the continuum variations in a simple manner. ",
        "introduction": "The Seyfert 1 galaxy \\mcg\\ has played an important role in studies of accretion onto black holes due to the presence of a broad, skewed iron line in its X-ray spectrum (Tanaka \\et 1995). The shape of the line seen with \\asca\\ is consistent with emission from the surface of an accretion disc extending from about 6 to more than 40 gravitational radii ($6-40r_{\\rm g}; r_{\\rm g}=GM/c^2$) inclined at about 30 deg to the line of sight (Fabian \\et 1995). Occasionally the red (i.e. lower energy) wing of the line is seen to extend below 4~keV (Iwasawa \\et 1996; 1999). This can be explained by the disc extending within $6r_{\\rm g}$ which may imply the black hole must be rapidly spinning. The presence of the broad iron line in \\mcg\\ has been confirmed by \\sax\\ (Guainazzi \\et 1999), \\xmm\\ (Wilms \\et 2001) and \\chandra\\ (Lee \\et 2002).  Here we present preliminary results from a long 325~ks observation of \\mcg\\ made with \\xmm. The source was at a similar flux level to the previous \\asca\\ observations, and about 70 per cent brighter than during the earlier 100~ks \\xmm\\ observation reported by Branduardi-Raymont \\et (2001) and Wilms \\et (2001). Simultaneous observations were made with \\sax, providing coverage from $\\sim$0.2--100~keV. The present work focuses on the spectrum above 2.5~keV and the iron-K line features; absorption and emission features below 2~keV due to oxygen and other elements will be discussed more fully in later work. ",
        "conclusions": "\\label{sect:disco}  A long observation with \\xmm\\ and \\sax\\ of \\mcg\\ in its typical state has again confirmed the presence of the broad, skewed iron line.  All the models considered above are formally acceptable ($\\chi_{\\nu}^{2}<1.0$), and all include emission within  $6r_{\\rm g}$, consistent with emission from a disc around a spinning black hole.  Model 4 provides the best (in a $\\chi^2$-sense) and most physically self-consistent explanation of the data.  In this model, the disc emissivity is described by a broken power-law in radius, and it is interesting to note that the break radius occurs at $\\sim 6r_{\\rm g}$. Beyond  this radius the disc has an  emissivity profile $q_{\\rm out} \\sim 2.5$ and produces an iron line with an equivalent width $\\sim 200$~eV (the $5.5-6.5$~keV core of the line shown in Fig.~\\ref{fig:fluxed_line}), both close to the values expected from standard accretion disc models. Within $6r_{\\rm g}$ the emissivity steepens, producing the strong low-energy tail to the line emission, also with an equivalent width $\\sim 200$~eV, suggesting additional physics within this region beyond that expected from standard accretion disc models (see Wilms \\et 2001). The strong iron line is consistent with an enhanced reflection spectrum and an overabundance of iron or ionisation of the disc surface. This last possibility will be examined in a later paper.  The iron line parameters do, however, depend slightly on the assumptions made about the complex absorption.  In particular all the above models formally assume the effects of ionised absorption are negligible above 2.5~keV, which is correct for absorption due to low-Z elements (e.g. O) but may not be true if there is absorption by higher-Z elements (e.g. Si, S). A preliminary analysis allowing for the possibility of Si and S edges did not change the requirement for a strong red wing to the iron line. A detailed analysis of the RGS data will yield constraints on the warm absorption (particularly on the low Z elements) and help remove existing degeneracies.  Future work to determine the spin must include emission from the immediate plunge region inside the innermost stable orbit (Reynolds \\& Begelman 1997; Agol \\& Krolik 2000) as well as returning radiation (Martocchia, Matt \\& Karas 2002)."
    },
    "0206/astro-ph0206185_arXiv.txt": {
        "abstract": "{ Key and still largely missing parameters for measuring the mass content and distribution of the Local Group are the proper motion vectors of its member galaxies. The problem when trying to derive the gravitational potential of the Local Group is that usually only radial velocities are known, and hence statistical approaches have to be used. The expected proper motions for galaxies within the Local Group, ranging from 20 to 100 $\\mu$as/yr, are detectable with VLBI using the phase-referencing technique. We present phase-referencing observations of bright masers in IC~10 and M33 with respect to background quasars. We observed the H$_2$O masers in IC10 three times over a period of two months to check the accuracy of the relative positions. The relative positions were obtained by modeling the interferometer phase data for the maser sources referenced to the background quasars. The model allowed for a relative position shift for the source and a single vertical atmospheric delay error in the correlator model for each antenna. The rms of the relative positions for the three observations is only 0.01 mas, which is approximately the expected position error due to thermal noise. Also, we present a method to measure the geometric distance to M33. This will allow re-calibration of the extragalactic distance scale based on Cepheids. The method is to measure the relative proper motions of two H$_2$O maser sources on opposite sides of M33. The measured angular rotation rate, coupled with other measurements of the inclination and rotation speed of the galaxy, yields a direct distance measurement. } ",
        "introduction": "An important astrophysical question is the nature and existence of dark matter in the Universe, which had been inferred originally from the flat rotation curves of galaxies (e.g. \\citeNP{FichTremaine1991}). The closest places to look for dark matter halos are the Milky Way and Andromeda galaxies in the Local Group. Various attempts have been made to weigh the galaxies in the Local Group and determine size and mass of the Milky Way and its not very prominent dark matter halo (\\citeNP{KulessaLynden-Bell1992}; \\citeNP{Kochanek1996}). Other attempts use Local Group dynamics in combination with MACHO data to constrain the universal baryonic fraction (\\citeNP{SteigmanTkachev1999}).  The problem when tyring to derive the gravitational potential of the Local Group is that usually only radial velocities are known and hence statistical approaches have to be used. \\citeN{KulessaLynden-Bell1992} introduced a maximum likelihood method which requires only the line-of-sight velocities (\\citeNP{HartwickSargent1978}), but it is also based on some assumptions (eccentricities, equipartition).  Clearly, the most reliable way of deriving masses is using orbits, which require the knowledge of three-dimensional velocity vectors obtained from measurements of proper motions. The usefulness of proper motions was impressively demonstrated for the Galactic Center where the presence of a dark mass concentration (presumably a black hole, see \\citeNP{MeliaFalcke2001}) has been umambiguously demonstrated by stellar proper motion measurements (\\citeNP{EckartGenzel1996}; \\citeNP{GhezKleinMorris1998}).  However, measuring proper motions of members of the Local Group to determine its mass is difficult. For the LMC \\citeN{JonesKlemolaLin1994} claim a proper motion of $1.2\\pm0.28$ mas/yr obtained from comparing photographic plates over a timespan of 14 years. \\citeN{SchweitzerCudworthMajewski1995} claim $0.56\\pm0.25$ mas/yr for the Sculptor dwarf spheroidal galaxy from plates spanning 50 years in time. \\citeN{Kochanek1996} shows that inclusion of these marginal proper motions can already significantly improve the estimate for the mass of the Milky Way, since it reduces the ambiguity caused by Leo I, which can be treated as either bound or unbound to the Milky Way. The same work also concludes that if the claimed optical proper motions are true, the models also predict a relatively large tangential velocity of the other statellites of the Galaxy. The dynamics of nearby galaxies are also important to determine the solar motion with respect to the Local Group to help define a standard inertial reference frame.  Despite the promising start, the disadvantage of the available optical work is obvious: a further improvement and confirmation of these measurements requires an additional large time span of many decades and will still be limited to only the closest companions of the Milky Way.  \\subsection{Proper Motions with the VLBA}  On the other hand, the expected proper motions for galaxies within the Local Group, ranging from 1 mas/yr to 20 $\\mu$as/yr, are relatively easy to see with VLBI using the phase-referencing technique. A good reference point is the motion of Sgr A* across the sky at a speed of 6 mas/yr reflecting the Sun's rotation around the Galactic Center at a speed of about 220 km/s. This motion is well detected between epochs separated by only one month with the VLBA (\\citeNP{ReidReadheadVermeulen1999}).  With the accuracy obtainable with VLBI one could in principle measure very accurate proper motions for most Local Group members within less than a decade. The main problem so far is finding appropriate radio sources. Useful sources would be either compact radio cores or strong maser lines associated with star forming regions. Fortunately, in a few galaxies bright masers are already known. Hence the task that lies ahead of us, if we want to significantly improve the Local Group proper motion data and mass estimate, is to make phase-referencing observations with respect to background quasars of known Local Group galaxies with strong H$_2$O masers.  \\subsection{Useful Local Group galaxies} The most suitable candidates for such a VLBI phase-referencing experiment are the strong H$_2$O masers in IC 10 ($\\sim$ 10 Jy peak flux density in 0.5 km/s line, the brightest known extragalactic maser; \\citeNP{BeckerHenkelWilson1993}) and IC 133 in M33 ($\\sim2$ Jy, the first extragalactic maser discovered). Both masers have been observed successfully with VLBI~(e.g. \\citeNP{ArgonGreenhillMoran1994}, \\citeNP{GreenhillMoranReid1993}). Additional fainter masers also exist in M33 that could be used to extend and improve the studies (e.g. for constraining galactic rotation; see section~\\ref{sec5}).  The two galaxies belong to the brightest members of the Local Group and are thought to be associated with M31. Their line-of-sight velocities are $-344$~km/s and $-180$~km/s respectively and are located at a distance of about 800 kpc. In both cases a relatively bright phase-referencing source is known to exist within a degree.  In addition their galactic rotation is well known from HI observations. Consequently, M33 and IC 10 seem to be the best known targets for attempting to measure Local Group proper motions with the VLBA. ",
        "conclusions": "The first results of our observations of H$_2$O masers in IC~10 and M33 have demonstrated the feasibility of high-precision astrometry at the 10 $\\mu$as level. With this accuracy we expect a 5 $\\sigma$ detection of the proper motions of IC~10 and M33 within one year. The H$_2$O masers in M33 were also observed, but the data reduction has not been finished yet. The observational techniques are nearly the same for IC~10 and M33 and so we expect promising results also for M33.  A possible pitfall for such a project is that individual maser components could be short lived and lost with time. However, in M33 the stronger maser components are known to exist for now two decades. An additional problem could arise if the observed phase referencing sources show motion in a core-jet structure which is unresolved with the VLBA. This could lead to a small apparent shift in the position of the phase referencing source which would be misinterpreted as a motion of IC~10 or M33. Using two calibrator sources, we are however able to exclude such a bias.  With the second and third set of observations in January 2002 and presumably in October 2002 we therefore will be able for the first time to detect significant proper motions in the Local Group out to 800 kpc."
    },
    "0206/astro-ph0206466_arXiv.txt": {
        "abstract": "During the course of an investigation on the interaction of the radio galaxy 3C\\,129 and its ambient cluster gas, we found excess X-ray emission aligned with the northern radio jet.  The emission extends from the weak X-ray core of the host galaxy $\\approx~2.5^{\\prime\\prime}$ to the first resolved radio knot.  On a smaller scale, we have also detected a weak radio extension in the same position angle with the VLBA.  Although all the evidence suggests that Doppler favoritism augments the emission of the northern jet, it is unlikely that the excess X-ray emission is produced by inverse Compton emission.  We find many similarities between the 3C\\,129 X-ray jet and recent jet detections from Chandra data of low luminosity radio galaxies.  For most of these current detections synchrotron emission is the favored explanation for the observed X-rays. ",
        "introduction": "The radio galaxy 3C\\,129 is a low luminosity (FRI type) 'tailed radio galaxy' seen in projection towards the outer edge of the X-ray emission from the hot gas of a nearby cluster of galaxies (Leahy and Yin, 2000, Taylor et al. 2001).  Since the cluster lies at low galactic latitude towards the anti-center, it has not been well studied in the optical.  We obtained Chandra observations in order to study the interaction of the radio structures with the hot intra cluster medium (ICM) and that work will be presented elsewhere (an analysis of the ICM properties has been performed by Krawczynski 2002, and a paper on pressure balance is in preparation).  In this paper we report on faint X-ray emission detected from the core of the 3C\\,129 galaxy and from the inner 3 kpc of the northern radio jet.  We include the results of 'follow-up' observations with the VLBA \\footnote{The National Radio Astronomy Observatory is operated by Associated Universities, Inc., under contract with the National Science Foundation.} in sec.~\\ref{sec:radio}.  X-ray emission from radio jets presents us with the problem of identifying the emission process but once this process is determined, we can then obtain new constraints on physical parameters (Harris and Krawczynski, 2002).  With the introduction of the relativistic beaming model of Celotti (Celotti, Ghisellini, \\& Chiaberge, 2001) and Tavecchio (Tavecchio, et al. 2000), most X-ray emission from jets has been interpreted as indicating either synchrotron emission or inverse Compton scattering off the cosmic microwave background (CMB).  For 3C\\,129, we show that synchrotron emission is the probable process, as has been found for a number of other FRI radio galaxies (Worrall, Birkinshaw, \\& Hardcastle, 2001; Hardcastle, Birkinshaw, \\& Worral 2001).  The implications of the detected X-ray emission are discussed in sec.~\\ref{sec:disc}.   The redshift of the radio galaxy at the center of the cluster, 3C\\,129.1 is z=0.0208 (Spinrad, 1975) and we take this for our distance estimate of D$_L$=126~Mpc with H$_o$~=~50~km~s$^{-1}$~Mpc$^{-1}$ and q$_o$~=~0. One arcsec then corresponds to 0.60 kpc. ",
        "conclusions": "\\label{sec:disc}  Granted that we are dealing with few photons and thus an insecure morphology, we believe the evidence favors synchrotron emission for the observed X-rays.  Undoubtedly there are bulk relativistic velocities in the jet producing the observed intensity differences between the N and S jets, but with velocity vectors not too far from the plane of the sky, we see only mild boosting and the parameters for IC/CMB emission are completely at odds with all other evidence.  If the bulk of the detected X-ray emission is in fact upstream of the radio knot N2.3, it would simply indicate that an acceleration region capable of producing $\\gamma$ of order 10$^7$ to 10$^8$ would be followed by a more extensive acceleration region incapable of such high energies, but rather producing up to $\\gamma~\\approx~10^4$.  Even in weak fields of order 30 to 50 $\\mu$G, the half-life for electrons producing X-rays is so short that they could travel no more than 30pc from their acceleration region.  Thus the X-rays clearly demarcate that sort of acceleration region.      There are now several detections of X-ray emission from jets in FRI radio galaxies.  For M87 (Marshall et al. 2002) the radio, optical, and X-ray morphologies are quite close if not identical but upstream offsets of X-ray brightness peaks compared to those of the radio have been documented for 3C\\,66B (Hardcastle et al. 2001) and 3C\\,31 (Hardcastle et al. 2002).  For these sources the offsets are a few hundred pc, slightly smaller than our (uncertain) value of 480pc for 3C\\,129.  The situation in the jet of Cen A (Kraft et al. 2000) is confused with some features aligning well at radio and X-ray bands, but for others it is not always clear which features correspond at the other wavelength.  In Table~\\ref{tab:fr1} we give comparative values of size and luminosities for several FRI detections.  It can be seen that the 3C\\,129 parameters are quite consistent with the others.  While we cannot rule out thermal bremsstrahlung as the cause of the X-rays from 3C\\,129, it seems likely that as for the other FRI detections, synchrotron emission is the favored process."
    },
    "0206/astro-ph0206246_arXiv.txt": {
        "abstract": "{Lithium abundances determined by spectral synthesis from both the 6708{\\AA} resonance line and the 6104{\\AA} subordinate line are reported for 11 Pleiades late-G and early-K stars observed at the William Herschel Telescope. Firm detections of the weak subordinate line are found for four objects, marginal detections for four, and upper limits for the remaining three stars.  Some of these spectra were previously analysed by Russell (1996), where he reported that abundances derived from the 6104{\\AA} line were systematically higher than those obtained from the 6708{\\AA} line by 0.2-0.7\\,dex. He also reported a reduced spread in the 6104{\\AA} line abundances compared with those determined from the 6708{\\AA} feature. Using spectral synthesis we have re-analysed Russell's data, along with our own. Our results do not entirely support Russell's conclusions. We report a $\\sim$0.7\\,dex scatter in the abundances from 6708\\AA\\ and a scatter at least as large from the 6104\\AA\\ line. We find that this is partly explained by our inclusion of a nearby \\ion{Fe}{ii} line and careful modelling of damping wings in the strong metal lines close to the 6104\\AA\\ feature; neglect of these leads to overestimates of the Li abundance which are most severe in those objects with the weakest 6104\\AA\\ lines, thus reducing the abundance scatter.  We find a reasonable correlation between the 6104{\\AA} and 6708{\\AA} Li abundances, although four stars have 6104{\\AA}-determined abundances which are significantly larger than the 6708{\\AA}-determined values by up to 0.5\\,dex, suggesting problems with the homogeneous, 1-dimensional atmospheres being used. We show that these discrepancies can be explained, although probably not uniquely, by the presence of star spots with plausible coverage fractions. The addition of spots does not significantly reduce the apparent scatter in Li abundances, leaving open the possibility that at least some of the spread is caused by real star-to-star differences in pre-main sequence Li depletion. \\keywords {stars: abundances - stars: late type - stars: interiors - open clusters and associations:individual: Pleiades} }  \\titlerunning{\\ion{Li}{i} 6104\\AA\\ in the Pleiades} ",
        "introduction": "\\label{sec-intro} The observed spread of lithium abundances in stars of the young ($\\sim$100 Myr) zero-age main sequence (ZAMS) \\object{Pleiades} cluster has been puzzling observers for around 20 years. The spread appears at $T_{\\mathrm{eff}}$ $\\sim$ 5800\\,K and continues into cooler stars. The magnitude of the spread increases with decreasing mass and reaches a maximum (more than 1\\,dex) at 0.8M$_{\\odot}$ in the mid-K stars before possibly decreasing in the coolest objects (Jones et al. \\cite{jones96}). Duncan \\& Jones (\\cite{dj83}) first reported the scatter, and interpreted it as being due to a large ($\\sim$0.4~Gyr) age spread within the cluster stars, although this has since been shown to be unlikely (Soderblom et al. \\cite{s93a}). A scatter is also seen in other young clusters (e.g. Randich \\cite{randich01}), but diminishes by the age of the Hyades (Soderblom et al. \\cite{s95} -- although most Li data for Hyades K stars are upper limit estimates). Pasquini et al. (\\cite{pasquini97}) report evidence that the spread reappears with a dispersion of~$>$1\\,dex among the solar-type stars of the old ($\\sim$5 Gyr) \\object{M67} cluster.  Soderblom et al. (\\cite{s93a}) noted that there is a connection between rotation and Li abundance in these cool stars, in the sense that rapid rotators seem to display higher Li abundances. Jeffries (\\cite{j2000}) showed that {\\em slow} rotators can have {\\em either} high {\\em or} low abundances. In addition, there is a correlation with chromospheric activity, with more-active stars also showing more lithium. It is not clear whether this link is causal or coincidental: {\\em i.e.} whether star spots or other chromospheric effects give rise to an {\\em apparent} abundance spread (Stuik et al. \\cite{stuik97}; King et al. \\cite{king2000}), or if the spread is real and linked to activity through a physical mechanism connected with the rotation in these objects (Soderblom et al. \\cite{s93a}; Jones et al. \\cite{jones97}). For instance, non-standard effects such as rotation-driven mixing (e.g. Chaboyer et al. \\cite{chaboyer95}) or perhaps even metallicity differences between individual cluster stars could cause variations in star-to-star lithium depletion during the PMS phase.  Almost all Li abundance measurements in main sequence Pop. I and Pop. II stars have been made using the strong \\ion{Li}{i}~6708{\\AA} resonance line. This line is formed high in the atmosphere and samples a limited range of depths. It is plausible that it could be affected by chromospheric activity (Houdebeine \\& Doyle \\cite{hd95}) or other inhomogeneities such as star spots (Giampapa \\cite{gia84}), leading to erroneous Li abundances. Stuik et al. (\\cite{stuik97}) argue that the effects of the chromosphere are primarily communicated through photospheric stratification and ionizing radiation fields, and therefore by changes in the ionization balance. To first order, all the alkali lines should be similarly affected by the presence of a chromosphere. Indeed, a number of authors have noted that the \\ion{K}{i} line at 7699\\AA\\ also shows some evidence for a spread in strength at the same effective temperatures, although perhaps not as large as for \\ion{Li}{i}~6708\\AA\\ (Soderblom et al. \\cite{s93a}; Jeffries \\cite{j99a}; King et al. \\cite{king2000}).   Explaining the observed spread in Li abundances among such a group of co-eval stars is an important goal. Either the abundance spread is real, which would tell us that non-standard mixing processes (those additional to convection) can produce star-to-star differences, probably as a result of differing angular momenta, or there really is no abundance spread; in which case the observations would tell us that crude, one-component, plane-parallel atmospheres poorly describe conditions in young, convective stars where there might be additional turbulence or atmospheric inhomogeneities to contend with.  A crucial test of the reliability of the atmospheres used in Li abundance analyses would be to obtain the abundance using alternative lines. Unfortunately the available optical lines (the subordinate transitions at 6104{\\AA} and 8126{\\AA}) are very weak and blended with other strong metal features. Nevertheless, these lines sample different depths within the atmosphere and the lower levels of their transitions form the upper level for the resonance line. It is, therefore, of primary importance to determine Li abundances from these lines separately, and to compare them with abundances measured from the resonance line.  Russell (\\cite{russell96}) measured 6104 (hereafter `6708' and `6104' are taken to mean `the \\ion{Li}{i} resonance line at 6708{\\AA}' and `the \\ion{Li}{i} subordinate line at 6104{\\AA}' respectively) in the atmospheres of six Pleiades late-G/early-K stars, and compared the abundances to those measured from 6708. His results suggested that 6104 gave abundances which were systematically higher than those from 6708 by 0.2-0.7\\,dex. He also reported a significantly smaller Li abundance spread derived from 6104. If Russell's conclusions were correct, they would indicate that something is seriously wrong with the model atmospheres used in such abundance determinations. The combination of a reduction in spread from one line over the other, and an abundance discrepancy between the two lines would suggest that there are problems with the temperature structure of the atmospheres, possibly related to inhomogeneities such as spots or plages on the stellar surface. However, Mart\\'{\\i}n (\\cite{martin97}) claims that Russell has probably overestimated the strength of the subordinate line by a factor of 2 or 3, due to his use of an inappropriate Gaussian-fitting technique. If this is taken into account, Mart\\'{\\i}n believes that the abundances from the two lines might agree to within 0.2\\,dex; in this case the agreement between the lines would support the validity of the atmospheric analysis, confirm the abundance spread, and suggest that non-standard mixing in PMS stars is important.  Given this controversy, we have repeated Russell's experiment, collecting our own data for several Pleiades G and K stars. We also obtained the data (including calibration frames) taken by Russell in 1993. These data were extracted, calibrated and analysed in exactly the same way as our own. ",
        "conclusions": "\\label{sec-spots} The main aim of this study was to investigate whether the Li abundances determined from the subordinate line at 6104{\\AA} agreed with those derived from the 6708{\\AA} resonance line. Such an agreement would lend confidence to the atmospheric approximations and NLTE corrections used in Li studies, and might imply that the spread in abundance previously reported from the 6708 line alone was indicative of a physical mechanism which could produce star-to-star differences in Li abundance during the approach to the ZAMS.  In our analysis, we find a scatter of $\\sim$0.7\\,dex from {\\em both} 6708 and 6104 (although the 6104 spread could be larger still, due to upper-limit estimates for some of the objects). The scatter is not explained by the uncertainties considered for either line. We find a broad correlation between the abundances from the two lines although for some stars there is an indication that the 6104-determined abundance is still 0.2-0.5\\,dex higher. The reason our results do not entirely agree with those of Russell (1996), who found that the 6104-derived abundances were higher and showed less scatter than those from 6708, can be traced to: (a) The large 6708-derived abundance for Hz 2311 found by Russell, which we have not been able to confirm. This increases the scatter in Russell's 6708-derived abundances; and (b) Russell's neglect of the strong, non-Gaussian damping wings of lines that blend with 6104 and the neglect of a weak \\ion{Fe}{ii} line, both of which might have been interpreted as Li absorption.  Although the abundance scatter derived from both lines supports the claim that the abundance spread is real, the disagreement in abundance for some cases does still illustrate a deficiency in our understanding of the stellar atmospheres.  Due to their ease of computation, one-dimensional atmospheres have been used for most abundance determinations until recently. However, they do not allow for the existence of inhomogeneities such as star spots on stellar surfaces, which are certainly known to exist in young, magnetically active, late-type stars and particularly in the Pleiades G and K stars (e.g. Krishnamurthi et al. \\cite{krish98}).  Doppler imaging study of the Pleiades G-type star (T$_{\\mathrm{eff}}$~=~5845\\,K) \\object{Hz~314} (Rice \\& Strassmeier \\cite{rice01}) showed that the star had cool spots at or near the pole, and within the equatorial regions. The temperatures of the spots varied from 4400\\,K to 5400\\,K, with the average being $\\sim$4700\\,K.  Spot areas have been measured spectroscopically for several highly-active G and K stars by O'Neal et al. (\\cite{oneal98}), who find spot coverage of 3-56\\%, and temperature differences between spotted and `quiet' (unspotted) regions ranging from 750\\,K to 1900\\,K.  The presence of spots could alter our conclusions in two ways: they might reduce the spread in Li abundances significantly in a group of stars with differing spot coverage (this has previously been explored by Soderblom et al. \\cite{s93a} and Barrado y Navasc\\'ues et al. \\cite{barrado01}, who conclude that the effect is not large enough). However, 6104 and 6708 might be affected differently due to their differing curve of growth temperature sensitivities, so differing spot coverage might well explain why some stars have agreement between 6708- and 6104-derived abundances while others reveal a higher abundance from 6104.  In order to probe what effect star spots might be expected to have on the abundances determined in this paper, we used a two-component, one-dimensional atmospheric simulation (TCODAS) technique included in the {\\sc uclsyn} code.  In the TCODAS models we balanced the spot and unspotted (quiet) star temperatures and areas, so that the luminosity-weighted average $T_{\\rm eff}$ would remain the same, using \\begin{eqnarray} aT_{\\mathrm{eff}}^{4} = a_{\\mathrm{cool}}T_{\\mathrm{cool}}^{4} + a_{\\mathrm{star}}T_{\\mathrm{star}}^{4} \\label{eqn-tcool-teff} \\end{eqnarray} where $a_{\\mathrm{cool}}$ and $T_{\\mathrm{cool}}$ are the cool-region area and temperature, and $a_{\\mathrm{star}}$ and $T_{\\mathrm{star}}$ are the area and temperature for the rest of the star. The areas are normalised to unity, such that: \\begin{eqnarray} a_{\\mathrm{star}} + a_{\\mathrm{cool}} = 1 \\end{eqnarray} $\\Delta T$ (= $T_{\\mathrm{star}} - T_{\\mathrm{cool}}$) and $a_{\\mathrm{cool}}$ were fixed at a number of discrete points, then Eq.~\\ref{eqn-tcool-teff} was solved for the appropriate $T_{\\mathrm{eff}}$. We assume that the cool regions are uniformly distributed across the star to avoid difficulties in modelling line-profile asymmetries.  The relative-light ($LR$) contributions from the two components were flux weighted using \\begin{eqnarray} LR_{\\mathrm{cool}} = \\frac{a_{\\mathrm{cool}} \\times T_{\\mathrm{cool}}^4}{T_{\\mathrm{eff}}^4} \\end{eqnarray} and \\begin{eqnarray} LR_{\\mathrm{star}} = \\frac{a_{\\mathrm{star}} \\times T_{\\mathrm{star}}^4}{T_{\\mathrm{eff}}^4} \\end{eqnarray}  We generated synthetic (LTE) spectra (at 6104 and 6708) using two component models with a luminosity-weighted average $T_{\\mathrm{eff}}$ of 5250\\,K, spot coverage of between 10 and 50 percent of the visible hemisphere and a temperature difference of either 1000\\,K or 1500\\,K between spotted and unspotted components. Finally we fitted our two-component syntheses with one component models at both 6104 and 6708. The resulting abundances are plotted on Fig.~\\ref{fig-lte2nlte}. These suggest that the abundances for Hz~263, Hz~2284 and Pels~19 (and Hz~129) could potentially be explained by the presence of star spots on these objects with coverage fractions of 10-50 percent and $\\Delta T$ of 1000-1500\\,K.  Note that the results obtained from TCODAS are LTE abundances, since we do not have access to NLTE syntheses, and application of 1-D NLTE corrections to the results would most likely not be valid.  Is there any evidence to support the notion that that those stars with significant discrepancies between their 6708- and 6104-derived abundances are the most magnetically active and spotted? We have no direct information on this and we must recognize that spot coverage might change with time, so unless observations of magnetic-activity indicators are co-temporal, they should be treated cautiously. There are indications of spot activity in the line profiles of Hz~263, Hz~2284 and Hz~129 (see Figs.~\\ref{fig-spec74}, \\ref{fig-russspec} and \\ref{fig-4-7}. Several line cores are filled in, especially the \\ion{Ca}{i} lines at 6102\\AA\\ and 6718\\AA. It is just such signatures which are used to make doppler maps of starspot distributions on more rapidly-rotating stars. However, we see little evidence of this in Pels~19, and similar signs of spot activity are present in Hz~2311 and Hz~522 (our observation rather than Russell's), where the 6708- and 6104-derived Li abundances {\\em might} be consistent (although the constraints are not strong and discrepancies of 0.2-0.3 dex cannot be ruled out).  A more quantitative magnetic-activity indicator could be the X-ray luminosity. Unfortunately, measurements are not available for all our stars and X-ray luminosity could easily vary by factors of 2-3 over time. A better diagnostic is the strength of H$\\alpha$ measured from the same spectra as the Li abundances. Soderblom et al. (\\cite{s93b}) have investigated H$\\alpha$ emission in the Pleiades. They found that late-G and early-K stars exhibit H$\\alpha$ absorption with a chromospherically-filled core. The amount of filling correlates with the rotation rate and other activity indicators. The strength of the underlying photospheric absorption is a function of $T_{\\rm eff}$. If, however, we make the reasonable assumption that this intrinsic absorption is uniform across the narrow $T_{\\rm eff}$ range considered here, then we can use the EW of the H$\\alpha$ core (1\\AA\\ either side of the line centre) to rank our targets in order of magnetic activity. These H$\\alpha$ core EWs are listed in Table~\\ref{table-results}. All of the targets exhibit H$\\alpha$ absorption, so it is the stars with the smallest EWs that are the most active -- namely, Hz~263, Hz~129 and Hz~2106. Two of these have discrepant 6708- and 6104-derived abundances, whilst the third might not have. However, Hz~2284 is ranked as one of the least active stars by this method, yet it does show an abundance discrepancy from the two lines, along with signs of spot activity in its Ca\\,{\\sc i} line profiles.  In summary the supporting evidence for a direct relationship between star spots and a discrepancy in Li abundance measured from the resonance and subordinate lines is ambiguous at present. The most active stars do seem to exhibit a discrepancy but there are counter examples as well. It should be made clear that the `spot' areas and temperature differences are not rigorously determined, and so should be treated only as fitting factors for use with the models, and not as actual areas or temperatures. We also remind the reader that changes in EW (and therefore changes in derived abundance) due to spot coverage (and spot coverage itself) vary over time. Jeffries et al. (\\cite{j94}) showed that 6708 was modulated by star spots during a rotation cycle in \\object{BD+22 4409}, a young, active star. The line strength was observed to vary by 10$\\pm$3\\% (equivalent to an abundance variation of $\\sim$0.15\\,dex within a one-component model) over the 10-hour rotation period. Over longer time-scales, Jeffries (1999) reported no variability larger than 0.1\\,dex on one-year timescales in a group of Pleiades K stars, but considered that there might be variations of 0.2-0.3\\,dex over 10 years. Finally, we note that there are differences in the two observations of Hz~522 used in this work: the lines in the spectra obtained by Russell are stronger.  This is consistent with an increase in spot coverage between the two sets of observations, resulting in an {\\em apparent} abundance decrease of 0.06$\\pm$0.02\\,dex between 1993 and 1998, assuming a one-component atmospheric model. If this were typical of the magnitude of this effect, it is clear that only a small fraction of the Li abundance spread could be accounted for in this way. Star spots could however, account for the majority of the discrepancies between abundances derived from 6104 and 6708 which arise due to the use of single-component models.  The fact that we can obtain agreement between the 6708- and 6104-derived abundances with relatively minor and plausible changes to the atmospheric structure (namely the introduction of star spots) without altering the size of the scatter observed in both the 6708 and 6104 abundances might support the idea that this abundance scatter is real. However, it is still possible that this agreement is coincidental and major revisions to our understanding of these young stellar atmospheres will be required. A cautionary note should be that scatter has also been observed in the potassium abundances of Pleiades stars from the \\ion{K}{i}~7699\\AA\\ line, which is formed under similar conditions to 6708 (Soderblom et al. 1993; Stuik et al. 1997; Jeffries 1999; King et al. 2000). As there are no plausible physical mechanisms to produce potassium depletion in these objects, this result still points to a deficiency in the atmospheric modelling. King et al. point to chromospheric activity altering the ionisation balance in the line formation region as being culpable. 6104 is formed a little deeper in the atmosphere than 6708, although there is plenty of overlap. It is beyond the scope of this paper to investigate whether plausible changes in the ionisation balance might also be able bring the 6104- and 6708-determined abundances into agreement, whilst at the same time significantly reducing the apparent scatter in Li and/or K abundances, though such an investigation is now called for."
    },
    "0206/astro-ph0206070_arXiv.txt": {
        "abstract": "A light echo around SN  1993J was observed 8.2 years after explosion by a HST WFPC2 observation, adding to the small family of supernovae with light echoes. The light echo was formed by supernova light scattered from a dust sheet, which lies 220 parsecs away from the supernova, 50 parsecs thick along the line of sight, as inferred from radius and width of the light echo. The dust inferred from the light echo surface brightness is 1000 times denser than the intercloud dust. The graphite to silicate fraction can not be determined by our BVI photometric measurements, however, a pure graphite model can be excluded based on comparison with the data.  With future observations, it will be possible to measure the expansion rate of the light echo, from which an independent distance to M81 can be obtained. ",
        "introduction": "Shortly after the burst of Nova Persei 1901, a light echo was observed to emerge with superluminal expansion. It was later properly explained (Couderc 1939) to be the nova light scattered by dust nearby which reached us later than the unscattered photons due to light travel effects.  The possibility of observing such an effect with supernovae was later discussed by many authors (e.g. Zwicky 1940, Schaefer 1987).  More than its splendid appearance, a light echo also sheds light on the circumstellar and interstellar environments through which it passes and the most famous example is from SN 1987A in Large Magelanic Cloud.  Observations of re-emission from the rings around SN 1987A provide us the information on the geometry, distribution and composition of its circumstellar medium (e.g., Lundqvist et al 1991) , help us to infer the stellar wind history of its pregenitor, and even enable us to determine an accurate distance to SN 1987A (Panagia et al. 1991). Also, the monitoring of its scattered light echoes enables a three dimenional mapping of the structure of the interstellar medium in front of SN 1987A (Xu et al. 1995), which reveals aspects of dense clouds and superbubbles that are difficult to reveal by other means.  This phenomenon has also been found in distant galaxies. For example, SN 1991T, a luminous Type Ia supernova in NGC 4527, exhibited a nearly flat light curve more than 950 days after maximum light, and spectral features that, although were present in earlier spectra, were substantially narrower and blueshifted on a significantly bluer continuum (Schmidt et al. 1994). Schmidt et al. attributed these features to a light echo, which was later confirmed by HST FOC observations (Sparks et al.  1999). Similar photometric and spectroscopic behaviors in the late-time observations of SN 1998bu have led to suspicion of a light echo (Cappellaro et al. 2001), but it has yet to be confirmed by direct high-resoulution imaging.  Due to limited spatial resolution, direct imaging observations of supernova light echoes are possible only in our local supercluster, which makes these phenomena rare events. Thus far, only Nova Persei 1901, SN1987A, SN1991T and SN1998bu are reported to have associated light echoes in the literature. Here we report on a light echo around SN1993J, the fourth such event.  The supernova SN1993J exploded on March 28, 1993 in the spiral galaxy M81, and due to its proximity (3.6 Mpc), it has been observed and at wavebands from radio to $\\gamma$-ray regions. It began as a type II supernova, but changed later to type Ib at the nebular stage, and was classified as type IIb. A series of observations with the Very Large Baseline Interferometer revealed an expanding radio shell that was decelerating (Bartel et al.  1999), reflecting the interaction between the shock front and the circumstellar medium. Intensive photometric and spectroscopic observations also showed an infrared excess after day 50, which may be indictive of an infrared echo (Lewis et al. 1994).  In this paper we report an optical scattered light echo around SN 1993J that was dicovered in an HST WFPC2 observation. In $\\S$2, we discuss our observations showing the light echo, and another archive HST WFPC2 observation with a nondetection of such a light echo. Models for this light echo are discussed in $\\S$3, which give the geometry and dust properties. For the distance to SN 1993J, we use the distance to M81, i.e., 3.63 Mpc ($\\mu=27.8\\pm0.2$, Freedman et al. 1994). ",
        "conclusions": "We have presented the WFPC2 observation of a light echo around SN 1993J after 8.2 years of supernova explosion. This light echo is due to light scattered from a dust cloud about $220 pc$ in front of SN 1993J, about $50 pc$ thick and $60 pc$ wide, revealing the existence of a sheet of dust (and gas) in another galaxy.  The dust inferred from the light echo surface brightness is 1000 times denser than the intercloud dust. The graphite to silicate fraction can not be determined by our BVI photometric measurements, however, a pure graphite model can be excluded based on comparison with the data.  Aside from studying the geometric structure of the interstellar medium in other galaxies, a light echo can be used to determine the distance to the host galaxy. To accomplish this, one needs to measure the expansion rate of the light echo. With future observations, we should be able to determine this expansion rate and obtain a measurement of the distance independent of that obtained by using Cepheid variables."
    },
    "0206/astro-ph0206300_arXiv.txt": {
        "abstract": "We present  a fractal dust model of the Universe based on Mandelbrot's proposal to replace the standard Cosmological Principle by his Conditional Cosmological Principle within the framework of General Theory of Relativity. This model turns out to be free from the de-Vaucouleurs paradox and is consistent with the SNe1a observations. The expected galaxy count as a function of red-shift is obtained for this model. An interesting variation is a steady state version, which can account for an accelerating scale factor without any cosmological constant in the model. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206136_arXiv.txt": {
        "abstract": "{ We present the 970-1175 $\\rm \\AA$ spectral energy distributions (SEDs) of 12 starburst galaxies observed with the Far Ultraviolet Spectroscopic Explorer {\\it FUSE}. We take benefit of the high spectral resolution of {\\it FUSE} to estimate a continuum as much as possible unaffected by the interstellar lines. The continuum is  rather flat with, in few cases, a decrease at  $\\rm \\lambda <~1050 \\AA$, the amplitude of which being correlated with various indicators of the dust extinction.  The far-UV SEDs are compared with synthetic population models. The galaxies with almost no extinction have a SED consistent with an on-going star formation over some Myrs. We derive a mean dust attenuation law   in the wavelength range 965-1140 $\\rm \\AA$ by comparing the SED of obscured galaxies to an empirical dust-free SED. The extinction is nearly constant longward of 1040 $\\rm \\AA$ but rises at shorter wavelengths. We compare our results with other studies of the extinction for galaxies and stars in this wavelength range.\\\\ ",
        "introduction": "The spectral energy distribution (SED) of star-forming galaxies in the far-UV (900-2000 $\\rm \\AA$) is known to be determined by the stellar initial mass function, the recent star formation history and dust extinction. As a consequence, the far-UV observations of galaxies are of fundamental importance to know the very recent star formation history in the universe and to interpret the spectra of high-z galaxies. \\\\   Thanks to the spectral observations of the {\\it IUE} satellite together with wide field UV imagers (SCAP (Donas et al. \\cite{donas87}) , FOCA (Milliard et al. \\cite{milliard94}, UIT  (Stecher et al. \\cite{stecher97})  or FAUST (Deharveng et al. \\cite{deharveng94}) experiments), our knowledge of the extinction longward of 1200 $\\rm \\AA$ has been considerably improved during the last 10 years. Even, if we are far from a complete understanding of the interplay of dust and UV emitting sources within galaxies, empirical laws have been found which allow us to estimate the extinction in galaxies, especially in starbursting objects (e.g. Calzetti et al. \\cite{calzetti00} and references therein). These successes may be traced back to the fact that the UV energy distribution for young starbursts are very similar. They can be fitted by a power-law (Leitherer \\& Heckman \\cite{leitherer95}) and changes in the exponent of the power-law can be attributed to reddening (Calzetti et al. \\cite{calzetti94}, Meurer et al. \\cite{meurer95}).\\\\  Stellar population models show that these properties  slightly change in the  900 - 1200 $\\rm \\AA$ domain. The time scale for reaching the equilibrium of UV flux in constant star formation gets shorter and, alternatively, age effects get more significant for instantaneous bursts. The SED cannot be fitted by a simple power-law, making the potential separation of age and reddening more difficult in practical terms. Observations of the spectral energy distribution of star-forming galaxies downward of 1200 $\\rm \\AA$ have been scarce so far and until recently were limited to those obtained with the Hopkins Ultraviolet Telescope (HUT): 19 spectra have been recently analyzed and an attenuation law for star-forming galaxies is derived from 900 to 1800 $\\rm \\AA$ by Leitherer et al. (\\cite{leitherer02}).  {\\it FUSE} ({\\it Far Ultraviolet Spectroscopic Explorer}) has recently opened again an access in the 900-1200 $\\rm \\AA$ range and the possibility of observing star-forming galaxies in this domain (e.g. Thuan et al. \\cite{thuan02}, Heckman et al. \\cite{heckman01a}). Although its very high spectral resolution is suited to the analysis of spectral lines,  {\\it FUSE}  can also be used to analyse the  continuum emission of starburst galaxies in an attempt to study their star formation history and their internal extinction. The high spectral resolution becomes an advantage for a better evaluation of the continuum in the presence of numerous absorption features as compared to the earlier work of Leitherer et al. (\\cite{leitherer02}) with HUT data. In the following, we report such an analysis for a sample of 12 starburst galaxies. ",
        "conclusions": "We  obtained the   spectral energy distribution of 12 nearby starburst galaxies from 970 to 1175 $\\rm \\AA$. The high spectral resolution of {\\it FUSE}  allowed us to properly estimate the continuum out of the interstellar lines. The general behavior is a rather flat distribution, in general agreement with the predictions of population synthesis models in this wavelength range for an active star formation. Nevertheless a  decrease of the flux downward 1050 $\\rm \\AA$ is observed in some objects: the mean ratio of the fluxes at 1070 and 1010 $\\rm \\AA$ is found to correlate with the metallicity and the extinction of the galaxies traced by their $\\rm F_{FIR}/F_{UV}$ ratio or the slope $\\beta$ of the spectral energy distribution longward of 1200 $\\rm \\AA$.\\\\ The galaxies   with an almost flat spectrum   have almost no extinction and a sub-solar metallicity (0.1 to 0.2$\\rm Z_\\odot$); their spectrum is well fitted by  a very young instantaneous burst of star formation or a constant  star formation and a Salpeter IMF. The FUV spectra of the galaxies which exhibit a decrease of the flux downward $\\rm 1050 \\AA$ cannot be accurately fitted by any of the models we consider. These galaxies show evidence for dust extinction and have in general a metallicity from $\\rm  0.4 Z_\\odot$ to $\\rm 2 Z_\\odot$. \\\\  We deduce a relative dust  attenuation law in the range 965-1140 $\\rm \\AA$ by comparing the SED of the sample galaxies with a dust free template constructed with the SED of the galaxies which do not exhibit any trace of extinction. The total amount of extinction is measured with the observed slope of the UV continuum longward of 1200 $\\rm \\AA$. We also derive an attenuation law related to the stellar color excess $A(\\lambda)/E(B-V)_s$ by using (and sligtly extrapolating) the calibrations of Calzetti et al. (\\cite{calzetti00}). We give a simple polynomial parametrization of $A(\\lambda)/E(B-V)_s$. The extinction that we find is almost constant from 1165 to 1040 $\\rm \\AA$ and rises at lower wavelengths. The agreement with the attenuation law proposed by Leitherer et al. (\\cite{leitherer02}) is good at $\\rm \\lambda > 1040 \\AA$ but we find a larger  extinction than Leitherer et al. at $\\rm \\lambda < 1040 \\AA$. Nevertheless, given the error bars in both studies,  the attenuation laws remain consistent.  \\begin{acknowledgement} This research is based on observations made with the NASA-CNES-CSA Far Ultraviolet Spectroscopic Explorer.  FUSE is operated for NASA by the Johns Hopkins University under NASA contract NAS5-32985.\\\\ This research has also 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}"
    },
    "0206/astro-ph0206408_arXiv.txt": {
        "abstract": "\\noindent {\\bf  The round and arc-shaped formations are known in some galaxies, the Bubble complex (the Hodge object) in NGC 6946 being the most remarkable. The rim of the complex has a form of a regular arc in which a part of a hexagonal structure is embedded. A similar morphology is recognized in  NGC 7421 galaxy, in that part of its rim which is leading in the galaxy motion through the intergalactic gas, as suggested by the bow-shock appearance of the HI halo of the galaxy. A hexagonal shape is also found in NGC4676A galaxy. The HII radial velocities across the Bubble complex  are compatible with its retrograde rotation and drift, which are characteristic for solitary vortices known in nonlinear hydrodynamics.  The drift motion may explain the location of the Bubble complex at the tip of the largest elliptical HI hole in NGC 6946. The hexagonal vortices in the atmospheres of Jupiter and Saturn are within the gas streams what seems to be suggestive as well. It is conjectured that the hexagonal rims of stellar systems might be relicts of flat segments of  the shock wave produced by the ram pressure. The giant stellar arcs in NGC 300 and M33 are associated with the energetic X-ray sources P42 and X-4 respectively (fig. 11), and these  might  be the relics of Hypernovae.}   \\vspace{0.3cm}  {\\bf Key words:} galaxies: individual (NGC 6946, M83, NGC 4449, NGC 7421, NGC 4676, LMC); stellar complexes; large scale star formation; spiral arms; vortices; ram pressure; Hypernovae ",
        "introduction": "Most stellar systems  concentrate to the center of mass and have more or less roundish  overall shape, without sharp rims (excepts the shell galaxies and some other products of the galaxy interactions). It is natural, for their shape is mostly determined by interplay of gravitation  and rotation (internal motions). However there are stellar complexes with sharp and very regular, arc-shaped rims, which  might have a peculiar origin  (Efremov 2001, 2002). Some  galaxies are known also to have sharp circular rims or the polygonal shape.  Here we  give examples of such geometrically regular systems and discuss hypotheses on their origin. ",
        "conclusions": "Putting the things together, we hypothesize that  the hexagonal appearance might be the transient  property of  persistent vortices (of wide size range),  which are suffered from the ram pressure. The straight segments of the rims might be a result of the the local flattening  of the shock  front on the space scale of local radius of curvature what  might, in turn, be due to stability of a flat shock wave against any weak perturbations that disturb its front surfaces (Chernin, 1999). If the star formation occurred in the shock wave while its front was flat, the respective part of the rim of the stellar system would preserve the hexagonal shape. The flatenning of the shock wave fronts seen in galactic spiral arms (Chernin 1999) and now in other formations worth to be investigated.  May be the Chandra data for clusters of galaxies will find something of the kind.  However, the origin of most  round or arc-shaped complexes is still a problem. Considering that both stellar arcs in NGC 300 and  M33 are the only such features in these galaxies and associated with the  most unusual in the respective galaxy X-ray sources, both well isolated,  the chance coincidence seems to be improbable. We believe that these X-ray objects, P42 and X-4 (especially P42), will be proved to be the remnants of Hypernovae. The estimated young age of the M33 X-4 SNR (Okada et al.,  2001) is surely incompatible with our suggestion, yet the source should be studied in more details. Both these arcs are rather  irregular, unlike  other arcs discussed here.  The similar origin is possible also for the LMC4 arcs, yet there are no X-ray sources located at the suitable positions in respect to the arcs. The X-ray binaries  concentrate  mostly to North of the LMC4 supershell, near the NGC 1978 cluster. The origin of the LMC arcs in the result of the ram pressure to the surface of impacting clouds seems to be more probable, considering also  the similar opening angles of these one-sided arcs. The main problem is that there are at least 2 and probably 5 arcs close to each other; it is a difficulty for any hypothesis, as well as the different orientations and ages.  The Hodge object is enigmatic also, with a couple of concentric semiarcs, suggesting action of a central pressure  yet having nor an evident source, neither an age - space pattern (see Fig. 1 here and  Fig. 9 in Larsen et al., 2002). The vortex motions might confine the young stars within the round complex; however, the sharp arc of the Western rim suggests the one-sided ram pressure action. Note that the drift of the complex might explain the absence of the HI hole around it, what is very strange considering the complex contains the cluster most suitable to form a supershell. May be the understanding of these features will come from an unexpected side, such as the dark matter  presence, or something else like this...  One may wonder how frequent are the round peculiar formations in galaxies. We were able to find  a dozen only  (Efremov 2001). The Hodge object in NGC 6946 was the only result of the special searches for the features similar to the giant arcs in the LMC4 region (P.Hodge, private communication). The appearance of this complex under different resolutions and contrasts suggests that many similar features may not be remarked. Under the low resolution, and in the distant galaxies they are practically  star-like (the best examples are the round complexes in M51 and especially in NGC 1232, see Fig. 12), whereas under the low contrast (such as in the Sandage-Bedke atlas) the encompassing circular rim is unnoted.  This is also the case for the stellar arcs even so large as in the LMC4 region. The rather similar in size complex in M83 (Comeron 2001) includes also two giant arcs (Fig. 12), which are  seen only in the best resolution images (Efremov, 2001).  At any rate, these peculiar entities are curious and elegant, and their interpretation may have the far reaching implications. Some were suggested in this  paper."
    },
    "0206/astro-ph0206122_arXiv.txt": {
        "abstract": "We used {\\it Hubble Space Telescope} WFPC2 images to identify six early-type galaxies with surface-brightness profiles that {\\it decrease} inward over a limited range of radii near their centers. The implied luminosity density profiles of these galaxies have local minima interior to their core break radii.  NGC 3706 harbors a high surface brightness ring of starlight with radius $\\approx 20$ pc. Its central structure may be related to that in the double-nucleus galaxies M31 and NGC 4486B. NGC 4406 and NGC 6876 have nearly flat cores that on close inspection are centrally depressed.  Colors for both galaxies imply that this is not due to dust absorption. The surface brightness distributions of both galaxies are consistent with stellar tori that are more diffuse than the sharply defined system in NGC 3706. The remaining three galaxies are the brightest cluster galaxies in A260, A347, and A3574.  Color information is not available for these objects, but they strongly resemble NGC 4406 and NGC 6876 in their cores. The thin ring in NGC 3706 may have formed dissipatively. The five other galaxies resemble the endpoints of some simulations of the merging of two gas-free stellar systems, each harboring a massive nuclear black hole. In one version of this scenario, diffuse stellar tori are produced when stars initially bound to one black hole are tidally stripped away by the second black hole. Alternatively, some inward-decreasing surface-brightness profiles may reflect the ejection of stars from a core during the hardening of the binary black hole created during the merger. ",
        "introduction": "Early-type galaxies are brightest in their centers and fade into the background at large radii.  There is no shortage of parametric forms that describe this smooth progression, but all more or less presume that the density of stars reaches its maximum in the center and decreases monotonically outwards.  Over the last decade, {\\it Hubble Space Telescope (HST)} imaging has shown that galaxy centers nearly always have singular surface brightness profiles of the form $\\Sigma_*(r) \\sim r^{-\\gamma}$ \\citep{crane, k94, f94, l95}. Low luminosity early-type galaxies, in general, have brightness profiles that are nearly power laws over several decades in radius with $\\gamma \\sim 1$ into the {\\it HST} resolution limit; for galaxies in the Virgo cluster this corresponds to radii of only a few parsecs. In contrast, the most luminous early-type galaxies have cores, defined by where the outer power law ``breaks'' or transitions to a shallower inner cusp --- but even there, $\\gamma > 0.$ \\citet{l95} clearly showed that even ``core galaxies'' had central {\\it density} cusps, $\\rho_L\\propto r^{-\\Gamma},$ with $\\Gamma$ significantly greater than zero; ``power-law galaxies'' typically had $\\Gamma\\sim2.$ \\citet{g96} verified the \\citet{l95} conclusions, showing that non-parametric inversion of the surface brightness profiles ratified the existence of density cusps in nearly all early-type galaxies imaged with {\\it HST.}  Massive nuclear black holes may play a critical role in the origin and survival of core structure \\citep{f97}. This hypothesis is motivated by the dichotomy in the central structure of elliptical galaxies. The low luminosity but dense power-law galaxies will be cannibalized by the high luminosity but more diffuse core galaxies. The long-term survival of low-density cores in luminous galaxies appears to demand moderation of any mergers by the black holes; the cores should have been filled in long ago without a strong tidal field to disrupt the in-spiraling nuclei of the power-law galaxies. The theoretical work of \\citet{mnm} verifies the \\citet{f97} argument that the initial creation of a core galaxy results from the merger of two power-law galaxies, each harboring a massive nuclear black hole.  Black holes may also be required to explain the double nuclei of M31 \\citep{l93} and NGC 4486B \\citep{l96}, two notable exceptions to the rule that stellar density reaches its maximum at the geometric galaxy center. The origin of these double nuclei is unknown, but their equilibrium and stability is most easily understood as a consequence of the massive black holes believed to reside in their centers \\citep{km31,k4486b}. The two visible nuclei in each galaxy are probably not distinct stellar systems, but may arise from a torus of stars bound to the black hole in a Keplerian potential \\citep{trem}.  Again, such tori may result from a merger in which two massive black holes are brought together by dynamical friction \\citep{hbr}.  In this context, we have searched for other galaxies with unusual central structures that may shed additional light on the formation of the central structure in galaxies, with particular attention to early-type systems exhibiting a central minimum in surface brightness. We identify six systems culled from a large sample of galaxies imaged by {\\it HST} with starlight distributions that do not neatly fit into the schema that cores always have cusps with $\\gamma>0.$ ",
        "conclusions": "It is possible that all six galaxies identified in this paper harbor stellar tori superimposed on normal cores. Certainly, the bright and sharply defined feature in NGC 3706 is difficult to explain as anything else. It is tempting to include M31 and NGC 4486B in this class, since the \\citet{trem} model for M31 is based on a stellar torus, even though two brightness maxima (``double nuclei\") in both galaxies are more pronounced and are less symmetric than in the galaxies described here. The origin of such tori will be discussed further below, as well as the hypothesis that some of the ``tori'' may really be cores that have been partially evacuated.  In either case, a genuine minimum in the density profile implies significant rotation, triaxiality, or anisotropy. An inward decrease in density cannot occur in spherical galaxies whose phase-space distribution functions depend only on the energy.  Such functions are guaranteed to be solutions of the collisionless Boltzmann equation. The density in this case for a spherical system is given by \\begin{equation} \\rho=2^{3/2}\\pi\\int_0^\\psi f(\\epsilon)(\\psi-\\epsilon)^{1/2}d\\epsilon, \\end{equation} where $f\\ge0$ is the mass per unit volume of phase space, $\\epsilon\\equiv -E$, and $\\psi\\equiv-\\Phi$ (the quantities $\\epsilon$ (total energy) and $\\psi$ (gravitational potential) are defined so that we work with non-negative variables: $\\psi(r)$ is positive at all radii, and $\\epsilon$ is positive for all bound stars). We can use \\begin{equation} d\\rho/d\\psi=2^{1/2}\\pi\\int_0^\\psi f(\\epsilon)(\\psi-\\epsilon)^{-1/2}d\\epsilon \\end{equation} to evaluate the density gradient anywhere in the galaxy (note that the derivative with respect to the upper limit of the integral is zero because the integrand goes to zero). Since $\\psi$ is a function only of radius, the gravitational field and any density gradient are both radial, and the latter is \\begin{equation} {d\\rho \\over dr} = {d\\psi \\over dr} \\, {d\\rho \\over d\\psi}. \\end{equation} Since $d\\psi/dr=-GM(r)/r^2$ is always negative and $d\\rho / d\\psi$ is always positive [$f(\\epsilon)$ is non-negative for all $\\epsilon$], $d\\rho /dr$  is always negative in any spherical system in which the distribution function depends only on energy. Thus, the systems observed in this paper are guaranteed to be either nonspherical or anisotropic, and may be both. All the systems for which we have dynamical observations indeed appear to have significant central rotation.  The rotation amplitude in NGC 4406 is modest, but it is strong in M31, NGC 3706, and NGC 4486B.  We next discuss two possibilities for the formation of the toroidal stellar systems.  \\subsection{Star Formation and Nuclear Disks}  There are several elliptical galaxies or bulges that harbor central stellar disks. \\citet{k01} present a short list of examples and argue that these disks were created {\\it in situ} by accreted gas funneled into the center. \\citet{svdb} emphasize that the disks can be radially compact and of high surface brightness. One might then ask if the tori in the present systems result from {\\it in situ} star formation as well.  Of the six galaxies, the relatively cold and dense ring in NGC 3706 comes closest to resembling a stellar disk formed directly from a gas disk.  One possible scenario is that the overall core structure was formed by the cannibalization of a low-mass stellar system, which contained some amount of gas, by NGC 3706. As is discussed in $\\S$\\ref{sec:n3706}, the ring in NGC 3706 is embedded in a larger stellar system that itself rises above and is twisted from the envelope brightness profile.  This system would contain the pre-existing stars in the denser portions of the cannibalized object, while the abrupt transition to the bright ring, itself, would reflect stars formed by gas delivered to the center of NGC 3706 in the same merger event. The issues remain: (1) Why has the ring not been filled in? (2) Why has the ring and surrounding stellar system not settled to the midplane defined by the envelope? Presently, the maximum ring density falls well inside the Roche radius of the estimated nuclear black hole mass. Thus, star formation in the ring requires either that the black hole was initially less massive, or that the pre-existing gas was in the form of dense molecular clouds, or that the gas layer in the disk was thinner than the current stellar disk.  The tori in the five remaining galaxies, however, appear to be considerably more diffuse and are less suggestive of thin stellar disks than the ring in NGC 3706.  Further, the transition from the outer cores to the radii of the tori appears to be smooth and gradual --- there is no abrupt transition to a high surface brightness ring, as is seen on the minor axis of NGC 3706. While it may be likely that these remaining tori were created in a merger event, as well (as we argue below), there is presently no compelling evidence that suggests that gas infall, followed by star formation within the cores, played a role in their formation.  \\subsection{Mergers of Galactic Nuclei Containing Black Holes}  It is now believed that the orbital decay of the massive black holes added to a galaxy by mergers with other galaxies may largely determine the inner distribution of starlight in the merger remnant. We speculate that the unusual structures seen in the present sample (with the possible exception of NGC 3706) may be evidence of this process. As it happens, both the ``stellar torus added'' and ``evacuated core'' interpretations of the observations may be supported by this scenario.  \\citet{bbr} argued that galaxies harboring massive black holes should occasionally merge and sketched out the orbital decay of the binary black hole that would be created in such an event.  Among other predictions, they suggested that in the later stages of the binary's life, the two black holes would eject stars from the newly merged core as the binary hardened, creating a local minimum in stellar density at the few parsec scale. \\citet{ebi} studied this problem numerically and showed that the binary black hole would generate a core with a shallow cusp in the merger remnant. Intriguingly, \\citet{mak97} showed that local minima in stellar density occurred within the core in some simulations.  The preservation of cores during mergers essentially demands that the final stages of such events are dominated by nuclear black holes \\citep{f97}. The central density contrast between core and power-law galaxies is strong enough that even modest cannibalism of the latter by the former should have filled in the diffuse cores long ago.  The ``core within a core'' merger endpoint hypothesized by \\citet{k84} prior to {\\it HST} observations, however, has never been found.  The tidal field of the nuclear black hole solves this problem by shredding the incoming nucleus prior to its final delivery to the center.  That cores exist only in the more luminous non-rotating elliptical galaxies, where gaseous dissipation may be less important during the merging of their progenitors, bolsters this picture \\citep{f97}.  All six galaxies discussed here appear to be core galaxies. The strong central isophote twist in NGC 3706 and the kinematically decoupled core in NGC 4406 suggests that mergers have influenced the inner structure of both galaxies. The increase in isophote ellipticity in NGC 6876 and A347 with decreasing radius, but at radii larger than at of the putative stellar tori, may also suggest cannibalization of a pre-existing stellar system. The cannibalized system would be less luminous, but denser and dynamically colder --- the ellipticity increase would reflect stars from such a galaxy being preferentially deposited in the core of the more luminous galaxy.  One immediate concern is why evacuated cores are so rare if they are a natural consequence of the process (decaying binary binary holes) that is supposed to make all cores. It is noteworthy that the recent simulations of \\citet{mnm} only produced the $\\gamma>0$ cusps seen in ``normal'' core galaxies. The nearly constant density cores generated by \\citet{mak97}, let alone cores with dips in density, were not seen. Significantly, \\citet{mnm} always merged two power-law galaxies to test the initial formation of a core as advocated by \\citet{f97}, while \\citet{mak97} started with cores already present in the merging galaxies. In a merging hierarchy, the more luminous galaxies may experience multiple episodes of merging or cannibalism.  If a merger product is used as the input for subsequent mergers, then there is a variety of possible initial structural forms. Two core galaxies may merge, a core galaxy may cannibalize a power-law galaxy, and so on; allowing for the concentration of the pre-merger cores provides an additional variable. If the \\citet{mnm} simulations are correct for the original formation of a core galaxy, but the \\citet{mak97} simulations are correct that central evacuation can occur with appropriate initial conditions, then perhaps the rarity of observational examples of evacuated cores merely reflects the likelihood of these conditions.  As the black hole binary hardens, it will eject stars from the merged core to increasingly large distances.  Early in its life, however, when the separation of the two black holes is similar to the core break radius, tidally stripped stars originally bound to either hole may still linger within the core. \\citet{hbr} emphasize that under some circumstances, such as when one galaxy has a relatively high nuclear stellar density and arrives at the center of the other galaxy with a high impact parameter, stars stripped from the incoming galaxy may form a diffuse torus.  \\citet{zb} have also discussed the formation of a stellar torus as one black hole strips stars from the other. In this picture, we really are seeing diffuse stellar tori superimposed against the cores of the merger remnants. Ironically, \\citet{hbr} attempted to explain the thin nuclear disks (as discussed in the context of NGC 3706) as being created by this process, but noted that they could only generate diffuse tori, not the high aspect-ratio disks seen in systems like NGC 3115 or NGC 4594. Some of the present systems may be the realizations of the toroidal structures seen in the \\citet{hbr} simulations."
    },
    "0206/astro-ph0206314_arXiv.txt": {
        "abstract": "{ We report spectroscopic observations ($400 - 800\\ nm$, $R\\approx100$) of Earthshine in June, July and October 2001 from which normalised Earth albedo spectra have been derived. The resulting spectra clearly show the blue colour of the Earth due to Rayleigh diffusion in its atmosphere. They also show the signatures of oxygen, ozone and water vapour. We tried to extract from these spectra the signature of Earth vegetation. A variable signal (4 to $10 \\pm3\\%$) around $700\\ nm$ has been measured in the Earth albedo. It is interpreted as being due to the vegetation red edge, expected to be between 2 to $10\\%$ of the Earth albedo at $700\\ nm$, depending on models. We discuss the primary goal of the present observations: their application to the detection of vegetation-like biosignatures on extrasolar planets. ",
        "introduction": "The search for life on extrasolar planets has become a reasonable goal since the discovery of Earth-mass planets around a pulsar (\\cite{wolszczan_et_al92}) and Jupiter-mass planets around main-sequence stars (\\cite{udry_et_al01}). Although the detection of Earth-mass planets is not foreseen before space missions  (like COROT scheduled for 2004, \\cite{schneider_et_al98}), it is likely that a significant  proportion of main sequence stars have Earth-like companions in their habitable zone.  An important question is what type of biosignatures will unveil the possible presence of life on these planets.  Spectral signatures can be of two kinds. A first type consists of biological activity by-products, such as oxygen and its by-product ozone, in association with water vapour, methane and carbon dioxide (\\cite{lovelock75}, \\cite{owen80}, \\cite{angel_et_al86}). These biogenic molecules present attractive narrow molecular bands. This led in 1993 to the Darwin ESA project (\\cite{leger_et_al96}), followed by a similar NASA project, Terrestrial Planet Finder (TPF, \\cite{angel_97}, \\cite{beichman99}). But oxygen is not a universal by-product of biological activity as demonstrated by the existence of anoxygenic photosynthetic bacteria (\\cite{blankenship_et_al95}).  A second type of biosignature is provided by signs of stellar light transformation into biochemical energy, such as the planet surface colour from vegetation, whatever the bio-chemical details (\\cite{labeyrie99}). This must translates into the planet reflection spectrum by some characteristic spectral features. This signature is necessarily a more robust biomarker than any biogenic gas such as oxygen, since it is a general feature of any photosynthetic activity (here leaving aside chemotrophic biological activity). Unfortunately, it is often not as sharp as single molecular bands: although it is rather sharp for terrestrial vegetation at $\\approx700\\ nm$ (\\cite{clark99}, \\cite{coliolo_et_al00}, see Fig.\\ref{vegetation}), its wavelength structure can vary significantly among bacteria species and plants (\\cite{blankenship_et_al95}).  \\begin{figure} \\centering \\includegraphics[width=8.75cm]{vegetation.ps} \\caption{Reflectance spectra of photosynthetic (green) vegetation, non-photosynthetic (dry) and a soil (from \\cite{clark99}). The so-called vegetation red edge (VRE) is the green vegetation reflectance strong variation from $\\approx5\\%$ at $670\\ nm$ to $\\approx70\\%$ at $800\\ nm$.} \\label{vegetation} \\end{figure}  Before initiating a search for extrasolar vegetation, it is useful to test if terrestrial vegetation can be detected remotely. This seems possible as long as Earth is observed with a significant spatial resolution (\\cite{sagan_et_al93}), but is it still the case if Earth is observed as a single dot? A way to observe the Earth as a whole is to observe the Earthshine with the Moon acting like a remote diffuse reflector illuminated by our planet. It has been proposed for some time (\\cite{arci12}) to look for the vegetation colour in the Earthshine to use it as a reference for the search of chlorophyll on other planets, but up to now, Earthshine observations apparently did not have sufficient spectral resolution for that purpose (\\cite{tikhoff14}, \\cite{danjon28}, \\cite{goode_et_al01}). We present in Section \\ref{section_EA} normalised Earth albedo spectra showing several atmospheric signatures. We show in Section \\ref{section_SR} how the vegetation signature around $700\\ nm$ can be extracted from these spectra. ",
        "conclusions": "Although it seems that the  Earth's vegetation signature might be visible as a red edge at $700\\ nm$, it is difficult to measure  in the Earthshine for two  reasons. The first reason is related to its variable amplitude, induced by a variable cloud cover and Earth phase. The second reason is because it is hidden below strong atmospheric bands which need to be removed to access the surface reflectance including the vegetation signature.   For the Earth, our knowledge of different surface reflectivities (deserts, ocean, ice etc) help us to assign  the VRE of the Earthshine spectrum to terrestrial vegetation. For an exoplanet, a VRE-like index might be as difficult to measure as for the Earth due to variable cloud cover of the planet. Even if an extrasolar planet would give a clear VRE-like spectral signal, its use as a biosignature would raise some questions because:  1/ For several organisms (such as {\\it Rhodopseudonomas}, \\cite{blankenship_et_al95}) the ``red edge'' is not at $700\\ nm$, but at $1100\\ nm$.  2/ Some rocks, like schists, may have a similar spectral feature. For instance, spectra of Mars show a similar spectral feature at $3.5 \\mu$, which were erroneously interpreted as vegetation due to their similarity with lichen spectra (\\cite{sinton57}).  We nevertheless believe that, associated with the presence of water (and secondarily oxygen) and correlated with seasonal variations, a vegetation-like spectral feature would provide more insight than simply oxygen on the bio-processes possibly taking place on the planet. But since water, and thus clouds and rain, are essential for the growth of vegetation, extrasolar planets with a very low cloud cover and a corresponding high vegetation index are unlikely, more especially if the planet is seen pole-on, with a bright white polar cover. On the other hand, an extrasolar planet vegetation surface could be larger than on Earth (like during periods in the paleozoic and mesozoic eras on Earth for example).  One must also note that the measurement of an extrasolar planet VRE will not suffer from the intrinsic difficulty of the same measurement for the Earth through the Earthshine spectrum: The extrasolar planet albedo will simply be given by the ratio of spectra $planet/mother\\ star$. But a model of the exoplanet atmosphere is necessary to be able to remove the absorption bands that may partially hide the vegetation. The detection of a VRE index between 0 and 10\\% requires a photometric precision better than 3\\%. Exposure time to achieve this precision with Darwin/TPF on an Earth-like planet at $10\\ pc$ with a spectral resolution of 25 is of the order of $100\\ h$ based on recent simulations (\\cite{riaud_et_al02}).  Finally, the Earth albedo spectral variations study is of interest for global Earth observation. It might provide data on climate change, as broad-band measurements recently showed (\\cite{goode_et_al01}). We also think that the spectrum of Earthshine might be used for example to monitor the global ozone (with the Chappuis or Huggins bands).  During the submission of the paper, we have been informed of a similar work by \\cite{woolf_et_al02}."
    },
    "0206/astro-ph0206064_arXiv.txt": {
        "abstract": "CH$_3$CN (J=6-5) was observed with a resolution of 2'' toward W75N using the BIMA interferometer. Two continuum sources were detected at 3 mm, designated MM1 and MM2 in previous studies. Alignment of two mm continuum sources with the outflow axis from MM1 suggests that these continuum sources may be the result of the outflow interacting with the interstellar medium. MM1 is coincident with compact CH$_3$CN emission. CH$_3$CN was not detected toward MM2. The distribution of optical depth ($\\tau _L$) is derived. An excitation analysis was not done because of large line optical depths. ",
        "introduction": "Hot cores in molecular clouds have a characteristic diameter of $\\sim$ 0.1 pc, density $\\ge 10^6 cm^{-3}$ and temperature $\\sim$ 100 K. Those hot cores that are luminous in the far IR (L$_{FIR} >$ 10$^4$ L$_{\\odot}$) but have no free-free radio emission are thought to be precursors to ultra-compact (UC) HII regions (Kurtz et al, 2000 and references therein). They are found in regions of massive star formation, often offset from UC HII regions (Hunter et al. 2000). Hot cores are principally studied through their molecular line emission and mm-submm continuum emission. Hofner et al. (1996) observed the UC HII region complex G9.62+0.19 and found that CH$_3$CN emission is centered on the youngest, densest component in this complex. This component also drives an energetic molecular outflow (Hofner et al. 1996) and has weak ($<$ 1 mJy) radio continuum emission at cm wavelengths (Testi et al. 2000). Based on observations of G9.62+0.29, Hofner et al. (1996) postulated that CH$_3$CN emission might be a tracer of UC HII precursors.  Wilner et al. (2001) used CH$_3$CN (J=12-11) observations of W49N to estimate the timescale for hot core evolution; they concluded that hot cores, which preceed the UC HII stage, have lifetimes less than or equal to UC HII regions. Zhang et al. (1998) observed many molecular lines, including CH$_3$CN (J=8-7), toward W51 and concluded from line asymmetries that CH$_3$CN was tracing infalling material. These studies found that CH$_3$CN tends to be: 1) tightly confined ($\\sim$ 10$^4$ - 10$^5$ AU), 2) associated with regions of massive star formation and 3) optically thick in mm-wave transitions. Watt \\& Mundy (1999) observed CH$_3$CN (J=6-5) toward 4 massive star formation regions (2 detections, 2 nondetections). Based on emission morphology, excitation analysis and chemical models by Millar et al. (1997), Watt \\& Mundy (1999) concluded that the hot core near G34.26+0.15 traced by CH$_3$CN emission was probably externally heated. In contrast to typical results from the observations cited above, they found no significant dust emission at 3mm. Based on single-dish detection and interferometer non-detection, they also concluded that the CH$_3$CN emission toward G11.94-0.62 has an extent of $\\sim$ 50,000 AU.  A CH$_3$CN (J=12-11) survey was undertaken using the 10m Heinrich Hertz Submillimeter Telescope (SMT) toward 48 known massive star formation regions to determine the fraction with detectable CH$_3$CN emission (Pankonin et al. 2001). Forty-six percent of the regions surveyed were detected in CH$_3$CN emission. Because the angular resolution was $\\sim$30'', they were unable to determine if CH$_3$CN originates from UC HII regions or from neighboring sites of even younger precursors of UC HII regions. Pankonin et al. (2001) found beam-averaged rotation temperatures and column densities to be consistent with internal heating, presumably from embedded protostars or the ionizing stars of the UC HIIs. They argued that beam dilution allowed only lower limits for CH$_3$CN column densities.  An investigation in the lines of CH$_3$CN (J=6-5) using interferometric resolution toward W75N is presented here to better determine the properties of this source and to establish how important resolution effects may be in the analysis of CH$_3$CN line emission. W75N was selected because it is a relatively nearby region of massive star formation region with strong CH$_3$CN emission and there are few previous observations of the hot core in this object. ",
        "conclusions": "CH$_3$CN (J=6-5) was observed toward W75N using the BIMA interferometer with 2'' resolution. Five compact ($\\sim$ 10,000 AU) K-components were detected (K=0-4) toward the peak of MM1 . CH$_3$CN emission is elongated with a deconvolved major-minor axis ratio of 3.4. The major axis coincides with the alignment of three cm-continuum sources reported previously with 0.1'' resolution VLA observations. Thus, the elongation is probably caused by multiple sources lying approximately along a straight line projected on the plane of the sky, although we cannot rule out CH$_3$CN tracing out an oblate cloud core. A bipolar outflow was previously detected toward MM1 in CO(J=1-0, J=3-2) and cm-continuum. MM2 and MM3 (detected by Shepherd 2001) lie on opposite sides of MM1 along the outflow axis, indicating that they may be heated by shocks as the outflow interacts with the local interstellar medium.  Analysis of K=2 and 3 emission lines indicates that CH$_3$CN (J=6-5) is optically thick ($\\tau _{max} \\approx$ 8). Thus, a radiative transfer/statistical equilibrium model which encorporates source kinematics and temperature and density gradients will be required to estimate T$_{rot}$ and N$_{CH_3CN}$. The results of this study support the previously observed pattern that CH$_3$CN in massive star formation regions tends to be compact ($\\sim$ 10$^4$-10$^5$ AU) and optically thick in mm-wave transitions."
    },
    "0206/astro-ph0206252_arXiv.txt": {
        "abstract": "{ The upper Asymptotic Giant Branch (AGB) is populated with oxygen rich and carbon rich Long Period Variables (LPVs). These stars are essential contributors to the near-IR light of intermediate age stellar populations. Individual observed spectra of LPVs are so diverse that they cannot be used directly in the synthesis of galaxy spectra. In this paper, the library of individual spectra of Lan\\c{c}on \\& Wood (2000) is used to construct averages that can be incorporated conveniently in population synthesis work. The connection between such spectra and stellar evolution tracks is discussed. In order to select a sorting criterion and to define averaging bins for the LPV spectra, correlations between their spectrophotometric properties are reexamined. While optical properties and broad baseline colours such as (I-K) are well correlated, a large dispersion is observed when these indices are plotted against near-IR ones. This is partly due to the intrinsic width of the upper AGB, which is illustrated by locating each of the multiple observations of individual LPVs on the HR diagram. It is argued that broad baseline colour-temperatures are the most sensible sorting criteria. The properties of the resulting sequence of average spectra indeed vary regularly. We further address: (i) the bolometric corrections and temperature scales needed to associate a spectrum with a given point on a theoretical stellar evolution track (or isochrone), (ii) the simplifying assumptions that will be implicitely made when using the average spectra, (iii) potential biases in the sample of Lan\\c{c}on \\& Wood and their effects, (iv) the small contribution of LPVs to the interstellar hydrogen emission lines in galaxies. It is emphasized that an a posteriori calibration of the effective temperature scale remains necessary, until consistent models for the evolution, the pulsation and the spectral appearance of LPVs become available. We suggest a recipe for the use of the average spectra at various metallicities. ",
        "introduction": "\\label{intro.sec}  The prediction of the near-infrared (near-IR) spectra of star clusters and galaxies is subordonate to the existence of stellar spectral libraries with a complete coverage of the evolved stages of stellar evolution. Stellar libraries are also necessary in the construction of near-IR colour-magnitude or two-colour diagrams, which may include narrow-band filter indices or absorption line measurements. A shortcoming of all libraries previously used in population synthesis work (e.g. Kleinmann \\& Hall \\cite{KH86}, Terndrup et al. \\cite{TFW91}, Lan\\c{c}on \\& Rocca-Volmerange \\cite{LRV92}, Pickles \\cite{Pick98}) is the poor coverage of the upper Asymptotic Giant Branch (AGB). The upper AGB hosts oxygen-rich and carbon-rich Long Period Variables (LPVs)\\footnote{ In agreement with common usage, LPVs include both semi-regular and Mira-type variables; Miras are LPVs with large (optical) amplitudes, i.e. $\\delta$V$>$2.5\\,magnitudes (Kholopov et al. \\cite{GCVS85}, Lloyd Evans \\cite{Llo83}, Hughes \\& Wood \\cite{HW90}).}, whose specific spectral features carry the potential of revealing the presence of intermediate age populations (Lan\\c{c}on et al. \\cite{LMFS99}). Indeed, upper AGB stars contribute of the order of 50\\,\\% of the K band light over a range of stellar population ages between $\\sim 10^8$ and $\\sim 10^9$\\,years (Persson et al. \\cite{PACFM83}, Renzini \\& Buzzoni \\cite{RB86}, Frogel et al. \\cite{FMB90}, Ferraro et al. \\cite{FFTetal95}, Bressan et al. \\cite{BGS98}, Girardi \\& Bertelli \\cite{GB98}, Lan\\c{c}on \\cite{Lan98}, Maraston \\cite{Mara98}, Mouhcine \\& Lan\\c{c}on \\cite{ML02_models}).  Recently, Lan\\c{c}on \\& Wood (\\cite{LW00}; hereafter LW2000) published a library of stellar spectra of luminous cool stars that includes a large sample of instantaneous observations of LPVs. The data span the wavelength range where most of the light of cool stars is emitted (0.5 to 2.5\\,$\\mu$m) with a spectral resolving power of $\\sim 1100$ longward of 1\\,$\\mu$m, and, with about 100 spectra, provide the most complete data set available to date.  The range of spectrophotometric properties observed in LPVs is large (LW2000, Alvarez et al. \\cite{ALPW00}). Even the classification of the empirical spectra has many caveats. Sorting algorithms based on optical or on near-IR criteria, on broad band colours or on spectroscopic properties, on mean or on instantaneous data, often give different results. The most conspicuous features, such as the H$_2$O vapour absorption bands around 1.4\\,$\\mu$m and 1.9\\,$\\mu$m (in oxygen rich LPVs), depend on details of the atmospheric structure that change from one pulsation cycle to the next. Such changes will be smoothed out in the integrated light of (large) stellar populations.  The purpose of the present paper is to provide a library of {\\em mean} stellar spectra of luminous cool variable stars, {\\em suitable for direct use in combination with a population synthesis model}. The library complements more readily available data for static red giants or supergiants. Both oxygen rich and carbon rich AGB stars are considered. We attempt to thoroughly discuss the impact of technical/physical choices that have to be made at various steps in the preparation and the use of the averaged spectra. \\medskip  The construction of a library of mean spectra requires the preliminary construction of suitable averaging bins. In principle, one would like to compute energy weighted mean spectra for individual stars, based on spectra taken at various phases over several pulsation cycles and on the light curve. However, this requires an enormous amount of spectroscopic and photometric data. Even the LW2000 data are not sufficient. This paper suggests that colour temperature is the best single parameter to order LPV spectra. The sample of useful input spectra is briefly described in Sect.\\,\\ref{sample.sec}, where the main selection criteria (and thus potential selection biases) are also recalled. Sequences of average spectra are presented in Sect.\\,\\ref{bin_choice.sec}. The statistical properties of the data that justify our selected sorting criteria are also presented in that section. \\medskip  The discussion of the uncertainties in population synthesis predictions inherent to the chosen averaging procedure must be based on the understanding of two processes that tend to disperse the stars of even a coeval population over quite a broad area of the Hertzsprung-Russell (HR) diagram: the TP-AGB thermal pulse cycles and the LPV pulsation cycles. The theoretical stellar evolution tracks used in population synthesis calculations provide the time evolution of the position of a star in the HR diagram. The thermal pulses are now commonly accounted for in the tracks, but the effects of pulsation, on timescales of hundreds or thousands of days, are not included. One may think of these tracks as providing the evolution of the static parent stars of LPVs. Unfortunately, building non-linear pulsation models and theoretical spectra for LPVs is extremely complex, and current models cannot yet provide the fundamental parameters of the parent star on the basis of one or several instantaneous empirical spectra or colours.  How, then, should one connect the averaged spectra of the new library with locations along the evolutionary tracks? Practical aspects of this question are addressed in Sect.\\,\\ref{useit.sec}. The subsequent discussion addresses fundamental difficulties. In particular, it emphasizes the need for a posteriori calibration of the relation between colour temperatures and effective temperatures (i.e. the temperature scale of the spectra). Potential selection biases in the LW2000 data, and their effect, are also discussed in Sect.\\,\\ref{disc.sec}. Recipes for extensions of the population synthesis calculations to other metallicities than quasi-solar are provided in Sect.\\,\\ref{metal.sec}. A concluding summary is given in Sect.\\,\\ref{concl.sec}. \\medskip  The averaged spectra are available in digital form through CDS\\footnote{Centre de Donn\\'ees Astrophysiques de Strasbourg, {\\tt http://cdsarc.u-strasbg.fr/CDS.html}, VizieR service}. ",
        "conclusions": "\\label{concl.sec}  This paper provides convenient and representative data for stars on the upper asymptotic giant branch, to be used in the studies of cool stellar populations of galaxies.  We have used the data of LW2000 to construct sequences of average spectra for O-rich and C-rich LPVs. The O-rich sequence is based on (I-K), which is taken as a first order effective temperature indicator. Despite the large dispersion between (I-K) (or other temperature indicators involving optical data) and near-IR properties, a regular behaviour is observed along the average sequence. This would not have been the case if sorting had been based on near-IR indices. For C-rich stars, both a temperature sequence and a sequence based on the C/O ratio are presented. For stars with thick dust envelopes (OH/IR stars and their carbon rich equivalent), the use of reddened versions of the above data is suggested.  In studies of stellar populations, the average spectra will be used in connection with stellar evolution tracks or isochrones. The latter are expected to provide the distribution of stars in the HR diagram, but also the onset of LPV pulsation (i.e. when to use the average LPV spectra instead of static giant spectra), the transition between O-rich and C-rich atmospheres, and the transition from an optically visible object to a dust-enshrouded far-IR source. As both the effective temperatures provided by stellar evolution tracks for the thermally pulsing AGB and the temperatures estimated for the average spectra are (independently) model-dependent and uncertain, we suggest caution when using the values indicated in this paper. The effects of reasonable changes in the temperature scales have been illustrated by Lan\\c{c}on et al. (\\cite{LMFS99}). We recommend an a posteriori calibration of the relative temperature scales of the spectra and the evolutionary tracks, based on observed properties of intermediate age star clusters of known ages and metallicities.  Using only one parameter, the effective temperature, to characterize LPV spectra is clearly a first order approximation (even if this particular choice of a parameter is arguably the most appropriate one). Pulsation period and amplitude as well as the phase in the pulsation cycle are expected to produce systematic effects, although neither the LW2000 sample nor the currently available stellar models are sufficient to pin those down. The period and amplitude distributions in the LW2000 data used here are biased towards high values, when compared to the distribution of LPVs (including Miras and semi-regular variables) found in star counts in the Milky Way or the Magellanic Clouds. As only TP-AGB stars with relatively high main sequence masses will ever contribute much to the integrated light of stellar populations, the problem is not as severe there. In the practice of population synthesis calculations, the errors resulting from this period and amplitude distribution cannot be disentangled from errors in the assumed onset instants of semi-regular and Mira-type pulsation."
    },
    "0206/gr-qc0206059_arXiv.txt": {
        "abstract": "A model of three-body motion is developed which includes the effects of gravitational radiation reaction. The radiation reaction due to the emission of gravitational waves is the only post-Newtonian effect that is included here. For simplicity, all of the motion is taken to be planar. Two of the masses are viewed as a binary system and the third mass, whose motion will be a fixed orbit around the center-of-mass of the binary system, is viewed as a perturbation. This model aims to describe the motion of a relativistic binary pulsar that is perturbed by a third mass. Numerical integration of this simplified model reveals that given the right initial conditions and parameters one can see resonances. These $(m,n)$ resonances are defined by the resonance condition, $m\\omega= 2n\\Omega$, where $m$ and $n$ are relatively prime integers and $\\omega$ and $\\Omega$ are the angular frequencies of the binary orbit and third mass orbit (around the center-of-mass of the binary), respectively. The resonance condition consequently fixes a value for the semimajor axis of the binary orbit for the duration of the resonance; therefore, the binary energy remains constant on the average while its angular momentum changes during the resonance. \\\\ \\\\ \\textbf{Key words:} celestial mechanics, relativity, gravitational waves ",
        "introduction": "In 1975, Hulse and Taylor found through the timing observations of binary pulsar PSR B1913+16 that the semimajor axis decayed at nearly the rate predicted by general relativity for the emission of gravitational radiation. This has served as significant indirect evidence of gravitational radiation reaction force as the result of energy balance due to the fact that gravitational waves leave the system. It has also shown the importance of the pulsar in probing strong gravity. Because of the universal nature of gravity, a correct analysis of astrophysical systems must involve the gravitational interaction between all of the masses in the environment plus other non-Newtonian effects. One would expect that over time, evidence of neighboring masses would appear in the timing observations of binary pulsars.  This study puts forth a model of three bodies, two of which constitute a relativistic binary pulsar and the third is a perturbing mass. This investigation extends the work of Chicone, Mashhoon, and Retzloff (1996a,1996b,1997a,1997b), which originally dealt with a binary system perturbed by normally incident gravitational waves. As in these previous models, the model derived here will be searched for resonances. In the following analysis, an $(m,n)$ resonance will occur when the relation $m\\omega=2n\\Omega$ is satisfied, where $m$ and $n$ are relatively prime integers, $\\omega$ is the frequency of the binary system, and $\\Omega$ is the frequency of the orbit of the third mass around the center-of-mass of the binary. Such resonances are physically noteworthy because they correspond to events where the collapse of the semimajor axis of the binary halts. That is, on average it stays fixed.  The N-body problem in general relativity theory is rather complicated; therefore, to capture the main effects of three bodies plus gravitational radiation reaction one can start with a classical three-body system with gravitational radiation damping as a linear perturbation. That is, in this work all post-Newtonian (relativistic) corrections will be neglected except for gravitational radiation reaction that will be treated in the quadrupole approximation. For the sake of simplicity, the internal structure of the masses will be neglected so that one is in effect dealing with three Newtonian point masses.  Two astrophysical systems that may demonstrate measurable behavior based on this model are the relativistic binary pulsar near a massive third mass and a planet in an orbit around a binary pulsar. While this study will focus on the former case the latter proves to be a timely point of interest. Astronomers have in the last six years found more than 50 stars which include at least one planet. The present number of stars that are currently being observed for planets is 1000 (Cameron, 2001). Pulsars have also proven to be detectable sources that show evidence of extrasolar planets. Finding a pulsar planet, where the pulsar is also part of a binary system is a possibility. Microlensing technique in its extrasolar planet search has turned up a possible planet with two suns (Bennet et al., 1999). Another possibility would be the future detection of a binary pulsar in a globular cluster.  A binary pulsar system that is situated in an approximately spherical cluster would be effectively the same as a binary attracted to a third mass equal to the mass of the matter inside the binary's orbit.  The timing of radio pulses from pulsars has opened up a rich and exciting branch of astronomy and astrophysics. Analysis of pulsar signals over time offers insight into the gravitational environment such as the possible presence of a companion mass. Hence, pulsar astronomy offers a fruitful method of detection for the astrophysical processes discussed in this paper. On the frontiers of astronomy, interferometric gravitational wave detectors based on Earth such as LIGO, VIRGO, GEO, AND TAMA aim to detect gravitational waves from known inspiraling binary systems. The shortening orbital separation that would arise as the result of energy loss via gravitational wave emission would result in an increasing frequency in the orbit. The gravitational wave signal would have twice this frequency (Blanchet, 2002). Thus, the predicted effect of a binary system's capture into resonance would leave its imprint in terms of the detected gravitational waves. Namely, the interval when the distance between the members of the binary stays fixed on average would give a gravitational wave signal that on the average would be of near constant frequency. Furthermore, the next generation of planned space-based detectors like LISA will broaden the scope of gravitational wave detection.  The initial approach to this problem in Section 2 involves setting up the equations of motion for the three masses with the influence of gravitational radiation reaction included. The radiation reaction force is understood to be very small, so it is viewed as a perturbation. The model is then simplified in Section 3 to make it more amenable to analysis. The results of the numerical integration of the simplified model are presented in Section 4. Next, section 5 expands upon the numerical results with a discussion about resonance. Section 6 contains concluding remarks.  The nonlinear disposition of the equations of motion for the relative motion of the binary system requires analysis that is suited for nonlinear behavior. Chicone, Mashhoon, and Retzloff (1997b) developed such an averaging method, that elucidates details about the structure of the orbit of a nonlinear system , especially when a resonance occurs. The system developed and studied here is suited for further analysis, in particular the application of the method mentioned above. ",
        "conclusions": "The main thrust of the analysis in this paper has been to introduce and develop a model of a binary system that is under the influence of gravitational radiation reaction and perturbed by a third body. The classical case of three gravitating bodies is considered here when the damping force due to the emission of gravitational waves by the system is included in the Newtonian equations of motion. After the pertinent equations of motion were derived, numerical analysis was done to explore important dynamics in the system -- namely, resonance between the orbits of the relative and third body motions. This would optimistically be an astronomically observable effect.  An analytical approach that would resolve details about this nonlinear system would add to the confidence of the numerical results. What remains, to augment the results of this work, is to pursue an analytical solution. The previous work of Chicone, Mashhoon, and Retzloff (1996,1997) outlines a novel averaging approach that looks into orbits near resonance. It was applied to the case of a binary system perturbed by the emission and absorption of gravitational waves. The model presented in this paper along with numerical indication of the existence of resonances leads naturally to this analytical approach for the model developed here (Wardell, to be published).  \\vspace{.5cm}  \\begin{flushleft} \\textbf{\\Large{ACKNOWLEDGEMENTS}}\\\\ \\vspace{.4cm} I would like to thank B. Mashhoon for his indispensable help and guidance in this project. I would also like to thank B. DeFacio for generously offering his computer facilities. \\end{flushleft}  \\begin{flushleft} \\textbf{\\Large{REFERENCES}}\\\\ \\vspace{.4cm} Bennet D.P., et al., 1999, Nature, 402, 57\\\\ Blanchet, 2002, gr-gc/0202016, to appear in Living Reviews in Relativity\\\\ Cameron A.C.,2001, Physics World, 14, 1\\\\ Chicone C.,Mashhoon B.,Retzloff D.G., 1996a, Ann. Inst. Henri Poincar\\'{e},\\\\ \\hspace{.5cm} Phys. Th\\'{e}or., 64, 87\\\\ Chicone C.,Mashhoon B.,Retzloff D.G., 1996b, J. Math. Phys., 37, 3997\\\\ Chicone C.,Mashhoon B.,Retzloff D.G., 1997a, Class. Quantum Grav., 14, 699\\\\ Chicone C.,Mashhoon B.,Retzloff D.G., 1997b, Class. Quantum Grav., \\\\ \\hspace{.5cm}14, 1831\\\\ Chicone C.,Kopeikin S.,Mashhoon B.,Retzloff D.G., 2001, Phys. Lett. A, \\\\ \\hspace{.5cm} 285,17\\\\ Landau L.D., Lifshitz E.M. 1971, The Classical Theory of Fields (Oxford: \\\\ \\hspace{.5cm}Pergamon Press)\\\\ Lyne Andrew G., Graham-Smith Francis 1998, Pulsar Astronomy (Cambridge: \\\\ \\hspace{.5cm} Cambridge University Press)\\\\ Schott G.A. 1912, Electromagnetic Radiation (London and New York: \\\\ \\hspace{.5cm}Cambridge University Press)\\\\ Stairs I.H., et al.,1989,Astrophys. J.,505,352 \\end{flushleft}  \\appendix  \\renewcommand{\\thesection}{APPENDIX \\Alph{section}:} \\setcounter{equation}{0} \\renewcommand{\\theequation}{\\Alph{section}\\arabic{equation}}"
    },
    "0206/astro-ph0206228_arXiv.txt": {
        "abstract": "There is a unique solution of the planet and star parameters from a planet transit light curve with two or more transits if the planet has a circular orbit and the light curve is observed in a band pass where limb darkening is negligible. The existence of this unique solution is very useful for current planet transit surveys for several reasons. First, there is an analytic solution that allows a quick parameter estimate, in particular of $R_p$. Second, the stellar density can be uniquely derived from the transit light curve alone. The stellar density can be used to immediately rule out a giant star (and hence a much larger than planetary companion) and can also be used to put an upper limit on the stellar and planet radius even considering slightly evolved stars. Third, the presence of an additional fully blended star that contaminates an eclipsing system to mimic a planet transit can be largely ruled out from the transit light curve given a spectral type for the central star. Fourth, the period can be estimated from a single-transit light curve and a measured spectral type. All of these applications can be used to select the best planet transit candidates for mass determination by radial velocity follow-up. To use these applications in practice, the photometric precision and time sampling of the light curve must be high (better than 0.005~mag precision and 5~minute time sampling). ",
        "introduction": "Planet transit searches promise to be the next big step forward for extrasolar planet detection and characterization.  Every transiting planet discovered will have a measured radius, which will provide constraints on planet composition, evolution, and migration history. Together with radial velocity measurements, the absolute mass of every transiting planet will be determined.  Transiting planets can be discovered around distant stars and in a variety of environments. Due to their special geometry many follow-up observations of transiting planets are possible, such as atmosphere transmission spectroscopy (note the first extrasolar planet atmosphere detection by Charbonneau et al.\\ 2002), search for moons and rings (Brown et al.\\ 2001), and detection of oblateness and the corresponding constraint on rotation rate (Seager \\& Hui 2002). Although no planet candidates discovered by the transit method have yet been confirmed by mass measurements, many searches are currently ongoing. The OGLE-III planet search (Udalski et al.\\ 2002) has observed numerous high-precision transit light curves from objects with small radii, including several potential planets.  The EXPLORE search (Mall\\'en-Ornelas et al.\\ 2002; Yee et al.\\ in preparation) has four potential planet candidates based on both photometric light curves and follow-up radial velocity measurements (Mall\\'en-Ornelas et al., in preparation).  The Vulcan planet search (Borucki et al. 2001) has some published results on transit candidates that, with radial velocity measurements, were determined to be eclipsing binary stars (Jenkins, Caldwell, \\& Borucki 2002).  Follow-up mass determination by radial velocity measurements are needed for planet transit candidates because late M dwarfs ($M \\geq 80 M_J$), brown dwarfs ($13 M_J < M < 80 M_J$), and gas giant planets ($M \\leq 13 M_J$) are all of similar sizes. This is due to a coincidental balance between Coulomb forces (which cause $R \\sim M^{1/3}$) and electron degeneracy pressure (which causes $R \\sim M^{-1/3}$). A high yield of confirmed planets from a list of planet candidates considered for mass follow-up is important, especially for planet searches with faint stars (e.g., fainter than 13th magnitude) which require relatively long exposures on 8-m class telescopes (e.g., 20--30 mins per star for a single radial velocity measurement). Hence understanding the transit light curves before follow-up can be crucial if a given project has a large number of planet transit candidates.  An analytical solution is always worthwhile to understand the general properties of a given physical system in an intuitive form. This is the case even when numerical fits are in practice the best way to determine the system parameters. In the case of a planet transit light curve we found that there is a unique solution for the five parameters stellar mass $M_*$, stellar radius $R_*$, companion radius $R_p$, orbital distance $a$, and orbital inclination $i$, under some important assumptions. This unique solution has several interesting applications --- especially when the photometric precision and time sampling are high --- including selection of the best planet candidates for follow-up mass measurements. Selection of the best candidates is especially important if the tendency for planets with small orbital distances to have low mass (Zucker \\& Mazeh 2002) --- and hence low radial velocity amplitudes --- is generally true.  In this case, on average, more effort to detect the planet mass via radial velocity variations of the parent star will be needed since the primary star's radial velocity amplitude scales linearly with planet mass.  The unique solution to a light curve with two or more transits was first mentioned in Mall\\'en-Ornelas et al.\\ (2002) where an approximate set of equations and a short description were presented. Sackett (1995) briefly touches on the unique solution by outlining parameter derivation with a known stellar spectral type, including a mention of period determination from a single transit. We begin this paper by describing the assumptions necessary to determine the unique solution in \\S2. In \\S3, for the first time, we present both the general set of equations that describe a planet transit light curve and their analytic solution. We discuss the errors in the parameters, and hence the limiting photometric precision and time sampling needed for the applications outlined in this paper in \\S4. The complications of limb-darkening are discussed in \\S5. In \\S6 we present four interesting applications that are made possible by the unique solution to the planet transit light curve. \\S7 concludes this paper with a summary. ",
        "conclusions": "We have presented the equations that describe a light curve with two or more transits and have presented the unique solution for the impact parameter $b$, the ratio of the orbital distance to stellar radius $a/\\rstar$, and stellar density $\\rhostar$.  Furthermore, with the stellar mass-radius relation we can uniquely derive the five parameters $\\mstar$, $\\rstar$, $i$, $a$, $\\rp$. This unique solution is only possible under the assumptions listed in \\S2, most importantly that the light curve is from a single star (i.e.  not two or more fully blended stars), that the planet is dark and is in a circular orbit, and the light curve transits have flat bottoms (which can be obtained at red or longer wavelengths).  We have found: \\\\ $\\bullet$ A simple analytical solution that can be used to quickly estimate the planet-star parameters, most importantly $\\rhostar$ and $\\rp$; \\\\ $\\bullet$ The stellar density can be uniquely determined from the transit light curve alone. Fitting codes that solve for star-planet parameters will find a number of best fits for different combinations of $\\mstar$ and $\\rstar$---these best fits will have the same stellar density. \\\\ $\\bullet$ For noisy data, the impact parameter $b$, $\\rstar$ and $\\rp$ are underestimated and $\\rhostar$ is overestimated due to a non-linear one-to-one correspondence for a given $\\Delta F$ between $b$ and transit shape (as parameterized by $t_F/t_T$).  The existence of the unique solution  for the above parameters allows several interesting applications, including: \\\\ $\\bullet$ The stellar radius and the planet radius can be estimated based on $\\rhostar$;\\\\ $\\bullet$ The likelihood that a shallow transit is due to contamination from a fully blended star can be estimated---with box-shaped transits being the least likely to be contaminated by blends; \\\\ $\\bullet$ A comparison of $\\rhostar$ as determined from the transit light curve with $\\rhostar$ from a spectral type can help identify blended stars and hence indicate that shallow eclipses are not due to planet transits; \\\\ $\\bullet$ In the presence of a blend the actual eclipsing companion radius is always larger than the radius derived from the light curve; \\\\ $\\bullet$ The period from a single transit event can be estimated with a known spectral type. \\\\  For most of these applications time sampling of $\\delta t < 5$ minutes and photometric precision of $\\sigma < 0.005$~mag are needed (or more generally $\\sigma^2 \\times \\delta t \\lesssim 1.5 \\times 10^{-4}$\\,mag$^2$\\,min).  This time sampling and photometric precision is reachable with current planet transit surveys (e.g., Mall\\'en-Ornelas et al. 2002). The transit shape must be well defined by high photometric precision and high time sampling because most star-planet parameters depend on transit shape. Specifically a limiting factor is time sampling of ingress and egress and their start and end times to determine transit shape."
    },
    "0206/astro-ph0206472_arXiv.txt": {
        "abstract": "{\\it Chandra} observation of the central region of the A1060 cluster of galaxies resolved X-ray emission from two giant elliptical galaxies, NGC 3311 and NGC 3309. The emission from these galaxies consists of two components, namely the hot interstellar medium (ISM) and the low-mass X-ray binaries (LMXBs). We found the spatial extent of the ISM component was much smaller than that of stars for both galaxies, while the ratios of X-ray to optical blue-band luminosities were rather low but within the general scatter for elliptical galaxies. After subtracting the LMXB component, the ISM is shown to be in pressure balance with the intracluster medium of A1060 at the outer boundary of the ISM\\@. These results imply that the hot gas supplied from stellar mass loss is confined by the external pressure of the intracluster medium, with the thermal conduction likely to be suppressed.  The cD galaxy NGC 3311 does not exhibit the extended potential structure which is commonly seen in bright elliptical galaxies, and we discuss the possible evolution history of the very isothermal cluster A1060\\@. ",
        "introduction": "The X-ray emission from elliptical galaxies directly tells us physical conditions of the interstellar medium (ISM), such as total mass, thermal energy, and chemical composition. Also, the number of low-mass X-ray binaries (LMXBs) gives us a clue how actively supernova explosions have occurred in the past.  Systematic relation between the X-ray and the optical blue-band luminosities, called the $L_{X}-L_{B}$ relation, was found with the {\\it Einstein} observatory \\citep{can87} and has been used to study the origin of the X-ray emission.  The hard component above $\\sim 3$ keV mainly comes from the LMXBs, and its luminosity is almost proportional to the optical one. The softer ISM component ($\\lesssim 1$ keV) shows a large scatter in the $L_{X}-L_{B}$ relation.  \\citet{mat01} analyzed PSPC data of 52 early-type galaxies and found that the galaxies are categorized into two groups, one is an X-ray luminous group ($L_{ISM:X}>10^{41}$ erg s$^{-1}$) characterized by extended hot halos with a radius of a few times 10$r_{e}$, and the other is an X-ray faint group showing compact halos.  As shown clearly in the case of NGC 4636 \\citep{mat98}, a large part of the emission in the X-ray luminous galaxies comes from an extended ($r \\sim 300$ kpc) hot ISM, indicating that they have extended potential wells around the galaxies with a scale as large as groups of galaxies.  In many cluster of galaxies, the central regions exhibit strong X-ray emission from low temperature ($kT \\leq 1$ keV) gases, which have been interpreted either due to cooling flows or to ISM in central cD galaxies.  The recent results for the cooling flow phenomena from {\\it Chandra} and {\\it XMM} indicated that the simple cooling-flow picture does not hold in a straightforward manner, and that some heating sources must be working in the center \\citep{tam01,kaa01,pet01,mol01,sas02}. The {\\it ASCA} study of the Perseus cluster \\citep{eza01,fur01} revealed extended ($\\leq 500$ kpc) cool emission in the center, suggesting a potential structure larger than that of the cD galaxy NGC 1275.  The excess central emission of the Centaurus cluster shows that the intra-cluster medium (ICM) is not isothermal but requires at least a two-phase gas \\citep{ike99}.  \\citet{pao02} found that the X-ray surface brightness profile around NGC 1399 in the Fornax cluster had three components, i.e.\\ cooling flow region, galactic and cluster halos. The formation process of these hierarchical potential structures around a galaxy or at the center of a cluster of galaxies is not yet understood.  Interaction between the ISM and the ICM may  cause a significant effect on the galaxy evolution.  Elliptical galaxies in the regions of high galaxy density tend to be X-ray faint \\citep{was91}. Recently, \\citet{vik01} studied the elliptical galaxies NGC 4874 and NGC 4879 in the center of the Coma cluster and found that the sizes of the ISM are as small as 3 kpc. The good spatial resolution of the {\\it Chandra} observatory enables us to look into the interaction process around bright galaxies in a number of clusters of galaxies.  A1060 (Hydra I cluster, $z=0.0114$) is an X-ray bright cluster of galaxies and is considered to be the archetype relaxed system. The {\\it ASCA} and {\\it ROSAT} observations of the cluster \\citep{tam96, tam00} showed that the temperature of the ICM is constant at $kT=3.1^{+0.3}_{-0.5}$ keV with the X-ray luminosity $L_{\\rm X}=2\\times 10^{43}$ erg s$^{-1}$ in the 2--10 keV band. An upper limit for the flux of the cool component assuming $kT=1 $ keV is $6\\times10^{41}$ erg s$^{-1}$ in the energy band 0.5--3 keV\\@. The X-ray morphology of the cluster is symmetric, and the surface brightness profile is approximated by a single $\\beta$ model or a modified ``universal'' NFW model \\citep{tam00}.  There are two giant elliptical galaxies in the center of the cluster. NGC 3311 is the cD galaxy with $m_{V}=12.65$. An E3 galaxy NGC 3309 with $m_{V}=12.60$ lies at only $1'.7$ (22 kpc in projection) away from NGC 3311, i.e.\\ within the core radius of the cluster ($3'.9\\pm 0'.1$) obtained by the $\\beta$ model fit for the PSPC surface brightness \\citep{tam00}.  The redshifts of the two galaxies correspond to 3593 km s$^{-1}$ for NGC 3311 and 4075 km s$^{-1}$ for NGC 3309, respectively \\citep{RC3}.  Assuming that NGC 3311 is settled at the center of the cluster potential as the cD galaxy, the velocity of NGC 3309 in the ICM would be at least 500 km s$^{-1}$. The apparently smooth and isothermal ICM and the presence of two giant galaxies makes this system a suitable object for the study of interaction features between galaxies and the ICM\\@.  We use  $H_{0}=75$ km s$^{-1}$Mpc$^{-1}$ and $q_{0}=0.5$, where the luminosity distance to A1060 is 46 Mpc and an angular size of $1''$ corresponds to 0.217 kpc. The solar number abundance of Fe relative to H is taken to be $4.68 \\times 10^{-5}$ \\citep{and89} throughout the paper. ",
        "conclusions": "{\\it Chandra} observation of the central region of the relaxed cluster of galaxies, A1060, confirmed that the temperature of the ICM was almost constant from outside to the inner 20 kpc region. The X-ray emission from the ISM of the two giant elliptical galaxies, NGC 3311 and NGC 3309, were clearly resolved. X-ray properties of the two galaxies are very similar and the ISM morphologies do not show a stripping feature even though the velocity difference between the two galaxies is $\\sim 500$ km s$^{-1}$. The ISM temperatures are constant at 0.7--0.9 keV, and the X-ray to optical luminosity ratios for the sum of the ISM and LMXB components are within the scatter for other galaxies.  The spatial extents of the ISM emission are, however, as small as $2 - 3$ kpc for both galaxies, and the pressure balance between the ISM and the ICM is achieved at the ISM boundaries.  We discuss, based on these results, that the ISM is mainly supplied from stellar mass loss and is confined by the external pressure of ICM, with highly suppressed heat conduction. The ICM in the central region of A1060 is relaxed, and the extended potential structure around the cD galaxies has not grown up yet, possibly due to gas mixing in past merger episodes."
    },
    "0206/astro-ph0206191_arXiv.txt": {
        "abstract": "{ We study a sample of 23 narrow-emission line galaxies (NELGs) which were selected by their strong variability as QSO candidates in the framework of a variability-and-proper motion QSO survey on digitised Schmidt plates. In previous work, we have shown that variability is an efficient method to find AGNs. The variability properties of the NELGs are however significantly different from those of the QSOs. The main aim of this paper is to clarify the nature of this variability and to estimate the fraction of AGN-dominated NELGs in this sample. New photometric and spectroscopic observations are presented, along with revised data from the photographic photometry. The originally measured high variability indices could not be confirmed. The diagnostic line-ratios of the NELG spectra are consistent with \\ion{H}{ii} region-like spectra. No AGN could be proved, yet we cannot rule out the existence of faint low-luminosity AGNs masked by \\ion{H}{ii} regions from intense star formation. ",
        "introduction": "The variability of flux densities is a common property of high-luminosity AGNs. We have performed a QSO search based on variability and proper motion (VPM survey) measured on a large number of digitised Schmidt plates in two fields (Meusinger et al. \\cite{Meu02}). The work in the M\\,92 field is the subject of the present series of papers. In the first paper (Brunzendorf \\& Meusinger \\cite{Bru01}; hereafter Paper\\,1), we discussed the motivation, the observational data, the data reduction procedure, and the selection of QSO candidates. The results from the follow-up spectroscopy and the properties of the resulting QSO sample were the subject of Paper\\,2 (Meusinger \\& Brunzendorf \\cite{Meu01}). The primary goal of the present study is to improve the understanding of the selection effects of this survey.  An object is considered a VPM-QSO candidate if it appears star-like, has no significant proper motion, and shows significant overall variability and long-term variability. The variability  is expressed by the indices $I_\\sigma$ (overall variability) and $I_\\Delta$ (long-term variability). For instance, an object with $I_\\sigma>2$ has a probability of $\\alpha>0.98$ to be variable. It is well known that high-luminosity AGNs vary on long timescales (years and longer). High priority QSO candidates have therefore to meet both $I_\\sigma\\ge2$ and $I_\\Delta\\ge2$. On the other hand, we found several QSOs with strong overall variability but without significant long-term variability (Paper\\,2). The long-term variability constraint may introduce a bias in the VPM QSO search, and it is therefore important to study also the subsample of variable, star-like objects with zero proper motion showing no significant long-term variability. In particular, we found 27 narrow emission line galaxies (NELGs) with redshifts $z \\la 0.2$ in this subsample. Most of these galaxies show strong emission lines. NELGs may be dominated by narrow-emission line AGNs (Seyfert\\,2, narrow-emission line Seyfert\\,1, LINERs), intense starbursts, or a mixture of both. For example, Ho et al. (\\cite{Ho97}) found that about half of the NELGs from their magnitude-limited sample show some form of AGN or composite spectra. In Paper~2, we have speculated that at least some of the VPM NELGs are dominated by AGNs, though the available data did not allow a clear-cut conclusion.  The present paper is concerned with the sample of the NELGs from the VPM survey.  The main question is whether the measured strong overall variability as well as the strong emission lines are related to AGNs or not. It is not our intention to provide a large and well-defined sample of NELGs useful for further detailed studies. Much larger samples (e.\\,g., Terlevich et al. \\cite{Ter91}; Ho et al. \\cite{Ho97}; Popescu \\& Hopp \\cite{Pop00}) are available and are better suited to the investigation of the overall NELG population. In Sect.\\,2, we present new spectroscopic and photometric observations. Section\\,3 is concerned with the variability properties of the NELGs. The spectroscopic properties are discussed in Sect.\\,4, and Sect.\\,5 reviews further properties of the galaxy sample. Sect. 6 concludes. As in the previous papers of this series, we adopt $H_0 = 50$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ and $q_0=0$. ",
        "conclusions": "We have studied the sample of NELGs from the VPM survey in the M\\,92  field. These objects have been selected  as QSO candidates because of their high variability indices (Paper~1) and have been classified later as NELGs (Paper~2). However, it was not clear from the previous data to what extent the variability and the spectral properties of the NELGs are related to AGNs. In the present paper, we re-investigated the variability and analysed the emission line-ratios, as well as the photometric and morphological properties of the NELGs. The main conclusions are the following. \\begin{itemize} \\item{The variability indices reported in Paper~2 primarily reflect increased measurement errors due to the resolved image profiles and do not provide evidence for AGNs. The measurement of variability for resolved objects requires techniques other than Gaussian profile fitting, even when the deviations from the stellar profiles are small. An unambiguous separation between stellar and nonstellar objects is crucial for the VPM survey. } \\item{The diagnostic line-ratio diagrams are best explained by \\ion{H}{ii} region-like spectra. None of the NELGs is unambiguously classified as an AGN. At least for some of the NELGs, the existence of LLAGNs  cannot be excluded. However, if present, AGNs do not dominate the integrated spectra. } \\item{The VPM NELGs are compact, blue galaxies. Most of them are related to starbursts. The sample consists of a range of various types, as is known from other samples of local starburst galaxies (e.\\,g., Coziol et al. \\cite{Coz98}). } \\end{itemize}  An important result of this work is the substantial improvement of both the photometric accuracy and the star-galaxy separation for the objects from the VPM survey in the M\\,92 field. This enabled us to identify additional VPM-QSO candidates with high or medium priority in our sample. Their spectroscopic observations, the newly detected QSOs and the properties of the enlarged QSO sample will be the subject of Paper~4 of this series."
    },
    "0206/astro-ph0206158_arXiv.txt": {
        "abstract": "Typical gamma-ray burst spectra are characterized by a spectral break, $E_p$, which for bright BATSE bursts is found to be narrowly clustered around 300 keV. Recently identified X-ray flashes, which may account for a significant portion of the whole GRB population, seem to extend the $E_p$ distribution to a broader range below 40 keV. Positive correlations among $E_p$ and some other observed parameters have been noticed. On the other hand, within the cosmological fireball model, the issues concerning the dominant energy ingredient of the fireball as well as the location of the GRB emission site are still unsettled, leading to several variants of the fireball model. Here we analyze these models within a unified framework, and critically reexamine the $E_p$ predictions in the various model variants. Attention is focused on the predictions of the narrowness of the $E_p$ distribution in different models, and the correlations among $E_p$ and some other measurable observables. These model properties may be tested against the current and upcoming GRB data, through which the nature of the fireball as well as the mechanism and site of the GRB emission will be identified. In view of the current data, various models are appraised through a simple Monte-Carlo simulation, and a tentative discussion about the possible nature of X-ray flashes is presented. ",
        "introduction": "\\label{sec:intro}  Gamma-ray burst (GRB) sub-MeV spectra contain most of the prompt information available from these mysterious sources. Progress in understanding the origin of such spectra, however, has been slow. A typical GRB spectrum is non-thermal, and can be fitted by a so-called Band-function (Band et al. 1993) with three parameters: a low energy power law photon index $\\alpha$, a high energy power law photon index $\\beta$, and the spectral break energy $E_p$ which defines the smooth transition between the two power laws. Since generally $\\beta < -2$, $E_p$ is related to the peak of the $\\nu F_\\nu$ spectrum, and therefore is also called ``E-peak\". Within the framework of the commonly considered cosmological fireball model, there are a number of variants (invoking different fireball contents, different emission sites, or different emission mechanisms) proposed to explain the GRB data. An important quantity which characterizes the GRB is $E_p$, which for 156 bright BATSE bursts in the 4B sample is found to be narrowly clustered around 300 keV, with a log-normal distribution with full width at half maximum of less than a decade (Preece et al. 1998; 2000). Theoretically, the value of $E_p$ is expected to be correlated with some other observational parameters, but such correlations could vary significantly in different models. These provide important criteria to evaluate the correctness of the models by comparing such correlations found in the data.  This topic has received heightened attention in view of two recent developments. On the one hand, a new category of X-ray transients, known as X-ray flashes (XRFs) has been identified (Heise et al. 2001; 2002, in preparation). The XRFs resemble normal GRBs in many respects, with the novelty that the peak energies are distributed from 100 keV to below 40 keV. These objects appear to form a natural extension of the GRB population in the softer and fainter regime (Kippen et al. 2001; 2002), which widens the narrow $E_p$ distribution found in the BATSE data, and it is estimated that XRFs represent a large portion (e.g. $\\siml 1/3$) of the whole GRB population. It is therefore interesting to see how present theories can accommodate this new category of GRBs. On the other hand, GRB light-curve variabilities (Fenimore \\& Ramirez-Ruiz 2000; Reichart et al. 2001) and spectral lags (Norris et al.  2000; Norris 2002) have been proposed as Cepheid-like luminosity indicators for the long duration ($t_b \\simg 2$ s) GRBs. These offer the exciting prospect of a direct relation between observable quantities and some of the theoretically most relevant parameters, such as the wind luminosity $L$, GRB intrinsic durations, emission-site magnetic fields, as well as the characteristic of the synchrotron or inverse Compton energies (which in many models are linked to the observed $E_p$). Empirically, some such correlations have been noticed. Lloyd-Ronning \\& Ramirez-Ruiz (2002) discovered a positive dependence of $E_p$ on the GRB variability (or luminosity). When combining this with the luminosity - variability correlation (Fenimore \\& Ramirez-Ruiz 2000; Reichart et al. 2001), one infers a positive correlation between $E_p$ and the burst luminosity $L$. Indeed, such a correlation has been seen in the BeppoSAX bursts with known redshifts (Amati et al. 2002). Thus, it is timely to critically revisit the physical $E_p$ predicted in various models, as a first step towards the goal of constraining or even identifying the nature of the fireball as well as the relevant emission site and mechanism for the GRB prompt emission. This is the purpose of the present paper. In \\S2.1, we present a synthesis of the current GRB model variants within a unified framework, and define the parameter regimes in which each model variant applies. In the rest of \\S2, we revisit these model variants, focusing specifically on the $E_p$ predictions. In \\S3, these predictions, as collected in Table \\ref{tbl-2}, are used to evaluate the current models through a simple Monte-Carlo simulation, and the possible nature of the X-ray flashes is tentatively discussed. Conclusions are drawn in \\S4. ",
        "conclusions": "\\label{sec:conc}  The nature of the GRB prompt $\\gamma$-ray emission, including the emission site and the energy contents of the fireball, are still poorly known after more than 30 years of efforts. We have analyzed the various fireball model variants within a unified picture, and have revisited the $E_p$ predictions of different models. These models are tested against the current GRB spectral data with a simple Monte Carlo simulation, with attention to the narrowness of the distribution and its dependence on some of the physical parameters. Our aim is to set up a general theoretical framework to allow unbiased tests of these models against the known data. Based on the analysis in this paper, we can evaluate the existing fireball models as follows.  1. The internal shock model is generally regarded as the most attractive candidate for the prompt $\\gamma$-ray emission of classical GRBs. The highly variable, spiky GRB lightcurves are naturally reproduced in such a model, and many studies have shown that it is successful in reproducing many of the GRB properties (e.g. Kobayashi, Piran \\& Sari 1997; Daigne \\& Mochkovitch 1998; Panaitescu, Spada, \\Mesz~ 1999; Spada, Panaitescu, \\Mesz~2000; Ramirez-Ruiz \\& Fenimore 2000; Guetta et al. 2001). Our findings in this paper indicate two important caveats for the internal shock model. First, unless the dispersions in the shock parameters (e.g. $\\epsilon_e$, $\\epsilon_B$, $\\theta_p$, $p$, etc) are very small or there exist some intrinsic correlations among the parameters, the internal shock model generates an $E_p$ distribution (Fig.1a) which, even in the optimistic case, is at least one order of magnitude broader than the straight (i.e. not otherwise corrected) distribution given by Preece et al. (2000). The calculated distribution may be still compatible with the data if XRFs are included in the observations, which intrinsically broaden the $E_p$ distribution\\footnote{It is worth mentioning that within the same burst, the internal shock model also tends to produce a wide $E_p$ distribution, which is incompatible with Preece et al. (2000) unless some fine tuning is made (S. Kobayashi, 2002, in preparation).}. Second, the synchrotron internal shock model may not be able to reproduce the $E_p-L$ positive dependence unless the $\\Gamma$ distribution is un-related or weakly related to $L$, or some further assumption is made (e.g. Ramirez-Ruiz \\& Lloyd-Ronning 2002).  2. The high-$\\sigma$ internal model could inherit most of the merits of the internal shock model, with some additional advantages such as a smaller $E_p$ dispersion and the right $E_p-L$ correlation. The caveat in such models is that they have so far been less well-studied than the internal shock models, including the basic particle acceleration and emission processes. It is also unclear how to lower the typical $E_p$ energy to the sub-MeV band, and the physical context of how a high-$\\sigma$ flow is launched in a collapsar is not well explored. More investigations in this direction are desirable (e.g. Blandford 2002).  3. In both the shock and magnetic dissipation internal models, both a narrow $E_p$ distribution and the right $E_p-L$ correlation are attainable if a pair photosphere is formed. However, it is unclear how common such a situation would be. The sharp lightcurve spikes may be also hard to reproduce.  4. The prompt $\\gamma$-ray emission predicted in external models is expected in most cases, unless the external medium is very under-dense or previous internal dissipation of the bulk kinetic energy has been highly efficient. Whether the observed GRB prompt emission is attributable to this component is in question. Our simulations show that these models produce too broad an $E_p$ distribution compared with other models, and probably a too steep $E_p-L$ correlation. Other arguments against such scenarios include the need for additional assumptions (blobs) and the inefficiency involved in interpreting the variability (Sari \\& Piran 1997, but see Dermer \\& Mitman 1999). An important feature of these models is that $E_p$ are positively dependent on the ambient ISM density $n_{ext}$, which is in principle a measurable parameter. This provides a potential test for the model. More broadband afterglow fits are needed before a significant statistical evaluation can be made. Even if the external scenario does not account for the GRB prompt emission, studies of such a model are nevertheless meaningful since they can explore the important bridge between the prompt emission and the afterglow.  5. The emission component coming directly from the baryon photosphere is expected to partially contribute to the observed GRB prompt emission. The $E_p$ distributions for different regimes are narrow, and $E_p-L$ correlations are easily accommodated at least in the regime where the emission is the strongest (i.e., $\\eta>\\eta_{c2}$). The best guess is that such a component appears mixed in with other components, and becomes important under certain conditions.  6. All the IC models tend to generate broader $E_p$ distributions as compared with their synchrotron counterparts. They are also less favored due to other reasons.  Finally, our entire discussion in this paper is within the context of a naked central engine. A more realistic scenario invokes the fireball-progenitor envelope interaction, e.g. in collapsar models, which is receiving increased attention. These would lead to additional emission components, which are beyond the scope of the current paper."
    },
    "0206/astro-ph0206335_arXiv.txt": {
        "abstract": "{ Diffuse X-ray emission from the Magellanic Clouds (MCs) is studied by using all the archival data of pointed ROSAT Position Sensitive Proportional Counter (PSPC) observations. For this purpose, contributions from the point and point-like sources in the ROSAT High Resolution Imager (HRI) and PSPC source catalogues are eliminated and periods of high solar activity are excluded. The spectral analysis yielded characteristic temperatures of $10^{6} - 10^{7}$~K for the hot thin plasma of the ISM which extends over the whole Large Magellanic Cloud (LMC) and the Small Magellanic Cloud (SMC). The total unabsorbed luminosity in the 0.1 -- 2.4~keV band within the observed area amounts to $3.2 \\times 10^{38}$\\,erg\\,s$^{-1}$ in the LMC and $1.1 \\times 10^{37}$\\,erg\\,s$^{-1}$ in the SMC, each with an uncertainty of $\\sim-40$\\%, $+100$\\%. The X-ray luminosity of the LMC is comparable to that of other nearby galaxies with pronounced star formation. In the LMC, hot regions were found especially around the supergiant shell (SGS) LMC\\,4 and in the field covering SGS LMC\\,2 and LMC\\,3. Highest temperatures for the SMC were derived in the southwestern part of the galaxy. The diffuse X-ray emission is most likely a superposition of the emission from the hot gas in the interior of shells and supershells as well as from the halo of these galaxies. ",
        "introduction": "The small distances to the Magellanic Clouds (MCs) allow us to separate the emission from distinct sources from that arising from surrounding gas within the galaxies. Not only the discrete point sources and extended supernova remnants (SNRs), but also diffuse emission coming from the interstellar medium (ISM) can be observed and studied in detail. In the interstellar space, stars are born out of the densest regions, whereas massive stars transfer matter back to the ISM in stellar winds and supernova explosions. Therefore, the diffuse component of the X-ray emission from galaxies will give us clues for a better understanding of the interaction between stars and the ISM as well as for the matter cycle within the galaxies.  With Einstein and ROSAT, high resolution X-ray imaging became possible and revealed diffuse X-ray emission in the Galaxy, the MCs, and other nearby galaxies \\citep[e.g.][]{1984ApJ...286..491F,1984ApJ...286..144W,1991Sci...252.1529S,1993namc.meet...59P}. It indicated the existence of a very hot component in the interstellar medium (ISM) with temperatures of $\\sim10^{6}$~K besides the cold gas observed in radio as well as warm components seen in the optical or UV. First supergiant shells (SGS) of relatively cold, ionized matter in the MCs were identified on H$\\alpha+$[\\ion{N}{ii}] images as filamentary structures \\citep{1978A&A....68..189G,1980MNRAS.192..365M}. Shells are interpreted to result from matter swept up by expanding gas, and can be also seen in \\ion{H}{i} emission maps as dense regions around voids in the neutral hydrogen distribution.  Based on Einstein IPC data, very hot gas in the Large Magellanic Cloud (LMC) was mapped by \\citet{1991ApJ...374..475W}, and further detailed analysis of the diffuse X-ray emission was performed \\citep[e.g.][]{1991ApJ...373..497W,1991ApJ...379..327W}. SGS LMC\\,2 \\citep{1980MNRAS.192..365M} is one of the supergiant shells discovered in the LMC and is located next to the star formation region 30 Doradus. \\citet{1991ApJ...379..327W} analyzed the Einstein IPC data of the SGS LMC\\,2 comparing it to infrared and \\ion{H}{i} observations, and found a ring of X-ray emission with a temperature of about $5 \\times 10^{6}$~K within a cavity in the \\ion{H}{i} map.  \\citet{1994A&A...283L..21B} analyzed ROSAT Position Sensitive Proportional Counter (PSPC) pointings covering the northern part of SGS LMC\\,4, which is located in the north of the LMC. They derived a temperature of $2.4 \\times 10^{6}$~K by fitting a thermal plasma spectrum. As for SGS LMC\\,2, ROSAT and ASCA observations made it possible to measure a plasma temperature of $k T \\approx 0.1 - 0.7$~keV \\citep{2000ApJ...545..827P}. Merged ROSAT PSPC images of the MCs were presented by \\citet{1994ApJ...436L.123S} and \\citet{1999IAUS..190...32S}. For the LMC, 140 PSPC pointings were used and 20 for the Small Magellanic Cloud (SMC), and images were created in different spectral bands of the detector. In the LMC, a hot plasma was found along the optical bar with temperatures between $4 \\times 10^{6}$ and $8 \\times 10^{6}$~K increasing from west to east. In the SMC no pronounced diffuse X-ray emission was detected.  Both the LMC and the SMC were observed by ROSAT \\citep{1982AdSpR...2..241T} in a period of over eight years in nearly 900 pointings, which covered the MCs almost completely. Therefore, a thorough study of the X-ray emission from the MCs has been started. The aim of this work was to establish a detailed picture of the high energy processes within a galaxy by producing a catalogue of sources in the MCs in the ROSAT band \\citep[see][]{1999A&AS..139..277H,2000A&AS..142...41H,2000A&AS..143..391S,2000A&AS..147...75S} and by analyzing the hot component of the interstellar medium. Since objects which can be detected in X-rays at the distances of the MCs are SNRs or binary systems including objects at the final stages of stellar evolution (supersoft sources, SSSs and X-ray binaries, XRBs), their distribution in combination with the structure and physical state of the ISM will indicate the region within the galaxies developing most actively and help us to understand the evolutionary history of the MCs. This work presents the spectral analysis of the complete ROSAT PSPC data unveiling the structure of the diffuse X-ray emission of the MCs with a spatial resolution of 15\\arcmin\\ $\\times$ 15\\arcmin\\,. ",
        "conclusions": "\\subsection{Temperature distribution}\\label{distrib_hot}  In order to verify the regions with significant diffuse X-ray emission and to look for correlations between the hot ionized gas and other components within the galaxies, the temperature distribution images (Fig.\\,\\ref{ktima}) were overlaid on other observations of the MCs, i.e.\\ merged image from PSPC data, DSS image, or \\ion{H}{i} column density maps. Fig.\\,\\ref{hikt_lmc} shows the \\ion{H}{i} column density map of the LMC with contours of X-ray temperatures, and in Fig.\\,\\ref{dsskt_smc}, the contours are plotted on a DSS image of the SMC.  \\begin{figure}[t] \\caption{\\label{hikt_lmc} Temperature contours from 0.1\\,keV to 0.4\\,keV in steps of 0.1\\,keV are superimposed on a \\ion{H}{i} map of the LMC \\citep{1998ApJ...503..674K} with positions of SNRs observed by ROSAT marked as squares.} \\end{figure}  \\begin{figure}[t] \\caption{ \\label{dsskt_smc} Overlay of the temperature distribution as contours from 0.1 to 0.8\\,keV in steps of 0.1\\,keV on a DSS image of the SMC. Boxes are SNRs observed by ROSAT and crosses are \\ion{H}{ii} regions.} \\end{figure}  \\subsubsection{LMC}  In the LMC, SGSs were found in \\ion{H}{i} and H$\\alpha$ observations with diameters of the order of 1\\,kpc, i.e.\\ each with a size of 5\\% -- 10\\% of the total size of the galaxy. The X-ray results show that there is an extended hot region in the eastern part of the LMC with a diameter of about 1\\,kpc, as can be seen in Fig.\\,\\ref{ktima}. This region covers the SGS LMC\\,2 \\citep{1980MNRAS.192..365M} where many active objects like SNRs and OB associations were found (see Figs.\\,\\ref{ktima}, \\ref{lumima}, and \\ref{hikt_lmc}). Using Einstein data, \\citet{1991ApJ...379..327W} derived a luminosity of $\\sim2 \\times 10^{37}$~erg~s$^{-1}$ for the diffuse emission from SGS LMC\\,2 assuming a temperature of $\\sim5 \\times 10^{6}$~K. Based on ROSAT PSPC data, \\citet{2000ApJ...545..827P} get a flux of $1.4 \\times 10^{-10}$~erg~cm$^{-2}$~s$^{-1}$ in the energy band of 0.44 to 2.04~keV, which corresponds to $L_{\\rm X} = 6 \\times 10^{37}$~erg~s$^{-1}$ (0.1 -- 2.4~keV). They used a model with $k T = 0.31$~keV. In the same region, we find $L_{\\rm X} = 1.7 \\times 10^{37}$~erg~s$^{-1}$ with a mean temperature of $k T_{\\rm MC} = 0.82$~keV.  High temperature was also determined at the northern rim of SGS LMC\\,4 and the boundary between SGS LMC\\,4 and SGS LMC\\,5. At such boundaries between the hot ISM and the cold dense supergiant shells, the heating mechanism is most effective, since the expanding gas hits on dense matter and is strongly decelerated, the kinetic energy transforming itself to thermal energy. For SGS LMC\\,4, the integration of the X-ray luminosity gives $L_{\\rm X} = 1.6 \\times 10^{37}$~erg~s$^{-1}$ and the mean value for the temperature is $k T_{\\rm MC} = 0.24$~keV. This result is in a very good agreement with the work by \\citet{1994A&A...283L..21B}, who obtained a temperature of $2.4 \\times 10^{6}$~K, i.e.\\ $k T = 0.21$~keV for the northern part of SGS LMC\\,4, as well as with the total luminosity of $\\sim 2 \\times 10^{37}$~erg~s$^{-1}$ for the SGS LMC\\,4 which was derived by \\citet{2000ApJ...545..827P} from the results of \\citet{1994A&A...283L..21B}, assuming that the northern emission is representative for the whole SGS LMC\\,4.  As for the total diffuse emission of the LMC, \\citet{1991ApJ...374..475W} determined a lower limit of $\\sim 2 \\times 10^{38}$~erg~s$^{-1}$ for the luminosity from Einstein data (0.16 -- 3.5~keV). The total luminosity derived from the ROSAT PSPC data in the energy band of 0.1 to 2.4~keV is $3.2 \\times 10^{38}$\\,erg\\,s$^{-1}$ (Sect.\\,\\ref{xlum36}) with an error of $\\sim-40$\\%, $+100$\\%, which is mainly due to uncertainties about the emission and absorption of the foreground gas.  \\subsubsection{SMC}  The hot gas in the SMC is well correlated with the optical main body which can be seen on a DSS image with superimposed temperature contours (Fig.\\,\\ref{dsskt_smc}). Regions with the highest temperatures are found in the southwest part of the galaxy with a large number of SNRs and \\ion{H}{ii} regions. The total X-ray emission not only arises from the optically visible part of the galaxy, but also shows the emission from the galactic halo. The existence of diffuse X-ray emission from the SMC was already reported by \\citet{1991ApJ...377L..85W} based on the analysis of Einstein data who obtained a total X-ray luminosity of the diffuse emission $L_{\\rm X} = 5.0 \\times 10^{38}$~erg~s$^{-1}$. But this was not verified in later ROSAT observations \\citep{1999IAUS..190...32S}. As it was shown in Sect.\\,\\ref{xlum36}, the ROSAT PSPC data yield  a much lower luminosity ($L_{\\rm X} = 1.1 \\times 10^{37}$\\,erg\\,s$^{-1}$). A possible explanation for the discrepancy by a factor of about 50 between $L_{\\rm X}$ derived from Einstein observation and that of ROSAT is the incompleteness of the list of the point sources which were removed from the Einstein data, since a large number of new point sources were found in subsequent X-ray observations.  Radio observations revealed shells and giant shells within the SMC and in the outer regions \\citep{1997MNRAS.289..225S}, but supergiant shells with large cavities like in the LMC are not known. The X-ray emission measured by the PSPC is the superposition of the emission from the interior of the shells which was heated up by stellar winds and supernovae.  \\subsection{Comparison to stellar distribution}  In the LMC, stellar associations forming large scale systems were found \\citep[Shapley's constellations,][]{1953PNAS...39..358M} suggesting secondary star formation \\citep{1983A&A...127..113B,1984A&A...131..347I}. The most prominent system is the Shapley's constellation III, located in the north of the LMC. It coincides spatially with the SGS LMC\\,4 and includes a large number of OB stars in about 20 young associations. The formation of these stars is thought to be caused by the gravitational instability of the SGS LMC\\,4. Another well known star formation region in the LMC is the 30 Doradus region located at the border between SGS LMC\\,2 and LMC\\,3. The comparison of the distribution of young stars in the LMC \\citep[$t_{\\rm age} \\leq 2 \\times 10^{7}$~yr,][]{1984A&A...131..347I} and the diffuse X-ray emission shows that the hot gas is well correlated with the distribution of young stars. In regions with highest temperatures, i.e.\\ around SGS LMC\\,4 or in the 30 Doradus region, there is a concentration of young supergiants ($t_{\\rm age} \\leq 8 \\times 10^{6}$~yr).  From optical observations it is known, that a young population of stars in the SMC is located along the bar \\citep[e.g.][]{1980A&A....87...92B}, and in particular concentrated in the southwest. In this part of the SMC, the hot gas detected in X-rays coincides with the most active regions. This can be verified in Fig.\\,\\ref{dsskt_smc} showing the positions of \\ion{H}{ii} regions, since associations of massive OB stars form \\ion{H}{ii} regions by their ionizing radiation and winds. Star forming regions are located in the eastern wing of the SMC as well, extending from the northern end to the east \\citep{1992MNRAS.257..195G}. Hot gas is distributed between the wing and the optical main body of the SMC, in a region around RA = 01$^{\\sl h}$ 04$^{\\sl m}$, Dec~=~--72\\degr~30\\arcmin\\ surrounded by SNRs and \\ion{H}{ii} regions in the north. The wing region is thought to be caused by dynamical interactions between the SMC and the LMC and/or between the MCs and the Galaxy which triggered star formation \\citep{1980PASJ...32..581M,1990A&ARv...2...29W}. A concentration of gas in the wing region was probably formed by forces acting between the halos of the MCs, rather than by an outflow of interstellar material from the main body of the SMC.  \\subsection{Diffuse emission from other nearby galaxies}  \\begin{table*}[t] \\caption{\\label{othergal} Properties of diffuse emission from some other nearby galaxies.} \\begin{tabular}{llccccp{5.7cm}l} \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{1} & \\multicolumn{1}{c}{2} & \\multicolumn{1}{c}{3} & \\multicolumn{1}{c}{4} & \\multicolumn{1}{c}{5} & \\multicolumn{1}{c}{6} & \\multicolumn{1}{c}{7} \\\\ \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{Name} & \\multicolumn{1}{c}{Type} & \\multicolumn{1}{c}{Distance} & \\multicolumn{1}{c}{\\ion{H}{i} mass} & \\multicolumn{1}{c}{N$_{\\rm H, Gal}$} & \\multicolumn{1}{c}{$L_{\\rm X}$} & \\multicolumn{1}{c}{References} \\\\ & & \\multicolumn{1}{c}{[Mpc]} & \\multicolumn{1}{c}{[$10^{09}$~M$_{\\sun}$]} & \\multicolumn{1}{c}{[$10^{20}$~cm$^{-2}$]} & \\multicolumn{1}{c}{[$10^{39}$~erg~s$^{-1}$]} & \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} SMC & SB(s)m & 0.06 & 0.4  & 5.1  & 0.01  & Stanimirovic et al.\\ (2000), this work \\\\ NGC 1705 & SA0 & 5.0 & 0.09  & 3.5  & 0.12  & Hensler et al.\\ (1998) \\\\ LMC & SB(s)m & 0.05 & 0.5  & 6.0  & 0.32 & McGee \\& Milton (1966), this work \\\\ NGC 1569 & IBm & 2.2 & 0.2  & 22  & 0.4  & Heckman et al. (1995) \\\\ NGC 4449 & IBm & 3.7 & 1.0  &  1.2  & 1.0  & Theis \\& Kohle (2001), Vogler \\& Pietsch (1997) \\\\ NGC 4631 & SB(s)d & 7.5 & 7.0  & 1.2  & 4.0  & Rand (1994), Vogler \\& Pietsch (1996) \\\\ NGC 253 & SAB(s)c & 2.6  & 1.0  & 1.3  & 4.0  & Puche et al.\\ (1991), Pietsch et al.\\ (2000) \\\\ NGC 4258 & SAB(s)bc & 6.4 & 5.0  & 1.2   & 20  & van Albada \\& Shane (1975), Vogler \\& Pietsch (1999) \\\\ M 83 & SAB(s)c & 8.9 & 20  & 4.0  &  36  & Huchtmeier \\& Bohnenstengel (1981), Ehle et al.\\ (1998) \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} ~\\\\*[2mm]Notes to\\\\ Col.\\ 2: Obtained from NASA/IPAC Extragalactic Database (NED).\\\\ Col.\\ 5: Galactic foreground \\nh\\ \\citep{1990ARA&A..28..215D}. \\\\ Col.\\ 6: Luminosity (0.1 -- 2.4~keV) of the diffuse emission from the galaxies corrected for Galactic foreground absorption.\\\\ \\end{table*}  As was shown in the last sections, the diffuse X-ray emission of the MCs seems to arise from the halo on the one hand, and from regions with high star formation activity on the other hand. This is in good agreement with the models for stellar bubbles and the evolution of stars, predicting that the ISM is heated up by stellar winds and supernova explosions. The hot gas in the ISM can blow out of the galactic disk and flow into the halo. Diffuse X-ray emission as evidence for the ISM heating is also observed in other nearby galaxies. Some examples are listed in Table \\ref{othergal} with distance, \\ion{H}{i} mass, foreground \\n{H}, and X-ray luminosity.  NGC 1704 and NGC 1569 are Magellanic-type nearby galaxies with ongoing star formation activities. Their \\ion{H}{i} mass is comparable to that of the MCs \\citep{1966AuJPh..19..343M,2000MNRAS.315..791S}. Similar to the MCs, there are bright star clusters embedded in superbubbles, and diffuse X-ray emission shows correlations with the H$\\alpha$ distribution. The unabsorbed X-ray luminosity of the diffuse emission is similar to that of the LMC \\citep{1998ApJ...502L..17H,1995ApJ...448...98H}. The face-on irregular galaxy NGC 4449 resembles the LMC as well, regarding its size, structure, and the diffuse X-ray emission \\citep{1997A&A...319..459V,2001A&A...370..365T}. For the outer disk and the halo, a temperature of $\\sim 3 \\times 10^{6}$~K was measured for the interstellar plasma, created probably during the star formation between $2 \\times 10^{6}$ and $4 \\times 10^{7}$ years ago.  Similarity in diffuse X-ray luminosity is also observed between the LMC and spiral galaxies, like the edge-on spiral galaxy NGC 4631 which is known to have many \\ion{H}{ii} regions and moderate star formation activity. The diffuse X-ray emission \\citep{1996A&A...311...35V} is about one order of magnitude higher than that of the LMC, as expected due to the bigger galaxy mass \\citep{1994A&A...285..833R}. Like the MCs, the gas distribution in NGC 4631 is thought to have been influenced by tidal interactions with its neighboring galaxies \\citep{1978A&A....65...47C}. In addition, supergiant \\ion{H}{i} shells like in the LMC were observed \\citep{1993AJ....105.2098R}.  One of the most famous nearby galaxies with X-ray emission studied in detail is the edge-on starburst galaxy NGC 253 with relatively low \\ion{H}{i} mass \\citep{1991AJ....101..456P}. \\citet{2000A&A...360...24P} analyzed the ROSAT data and found that the contributions of the nuclear area, disk, and halo to the diffuse X-ray emission are about equal. The nuclear area mainly consists of a heavily absorbed source with $k T = 1.2$~keV and $L_{\\rm X} = 3 \\times 10^{38}$~erg~s$^{-1}$ and an 'X-ray plume' described by two components ($k T = 1.2$~keV and $k T = 0.33$~keV), which is thought to originate from the interaction between the galactic wind from the starburst nucleus and the interstellar medium within the disk. The halo emission is very soft with a temperature of $k T \\approx 0.1$~keV. In the spectra of the MCs, this component is difficult to be verified because of the higher Galactic foreground absorption.  Further spiral galaxies are known to show diffuse X-ray emissions, like the Seyfert 2 galaxy NGC 4258 \\citep{1975A&A....42..433V,1999A&A...352...64V} or the face-on spiral galaxy M 83 \\citep{1981A&A...100...72H,1998A&A...329...39E}. These galaxies show at least two diffuse emission components, the soft halo emission and the hard, highly absorbed disk emission. Their diffuse X-ray luminosity is very high ($L_{\\rm X} > 10^{40}$~erg~s$^{-1}$) what is to be expected due to their star formation activity."
    },
    "0206/astro-ph0206103_arXiv.txt": {
        "abstract": "We combine weak lensing measurements from the Red-Sequence Cluster Survey (RCS) and the VIRMOS-DESCART survey, and present the first direct measurements of the bias parameter $b$ and the galaxy-mass cross-correlation coefficient $r$ on scales ranging from 0.2 to 9.3 $h_{50}^{-1}$ Mpc (which correspond to aperture radii of $1.5'-45'$) at a lens redshift $z\\simeq 0.35$. We find strong evidence that both $b$ and $r$ change with scale for our sample of lens galaxies ($19.5<R_C<21$), which have luminosities around $L_*$. For the currently favored cosmology $(\\Omega_m=0.3,~\\Omega_\\Lambda=0.7)$, we find $b=0.71^{+0.06}_{-0.04}$ (68\\% confidence) on a scale of $1-2h_{50}^{-1}$ Mpc, increasing to $\\sim 1$ on larger scales.  The value of $r$ has only minimal dependence on the assumed cosmology. The variation of $r$ with scale is very similar to that of $b$, and reaches a minimum value of $r\\sim 0.57^{+0.08}_{-0.07}$ (at $1 h_{50}^{-1}$ Mpc; 68\\% confidence). This suggests significant stochastic biasing and/or non-linear biasing. On scales larger than $\\sim 4 h_{50}^{-1}$ Mpc the value of $r$ is consistent with a value of $r=1$.  In addition we use RCS data alone to measure the ratio $b/r$ on scale ranging from 0.15 to 12.5 $h_{50}^{-1}$ Mpc ($1'-60'$) and find that the ratio varies somewhat with scale. We obtain an average value of $b/r=1.090\\pm0.035$, in good agreement with previous estimates.  A (future) careful comparison of our results with models of galaxy formation can provide unique constraints, as $r$ is linked intimately to the details of galaxy formation. ",
        "introduction": "The growth of structures in the universe via gravitational instability is an important ingredient in our understanding of galaxy formation. However, the connection to observations is not straightforward, as we need to understand the relation between the dark matter distribution and the galaxies themselves. Galaxy formation is a complex process, and it is not guaranteed a priori that this relation, referred to as galaxy biasing, is a simple one. The bias might be non-linear, scale dependent or stochastic. In the simplest case, linear deterministic biasing, the relation between the dark matter and the galaxies can be characterized by a single number $b$ (e.g., Kaiser 1987).  Most observational constraints of biasing come from dynamical studies (see Strauss \\& Willick 1995) which probe relatively large scales ($10 h_{50}^{-1}$~Mpc or more). Recent estimates on these scales suggest values of $b\\sim 1$ for $L_*$ galaxies (e.g., Peacock et al 2001; Verde et al. 2001). On smaller scales some constraints come from measurements of the galaxy two-point correlation function, which is compared to the (dark) matter correlation function computed from numerical simulations.  \\vbox{ \\vspace{0.5cm} \\footnotesize \\noindent $^1$~CITA, University of Toronto, Toronto, Ontario M5S 3H8, Canada\\\\ $^2$~Department of Astronomy and Astrophysics, University of Toronto, Toronto, Ontario M5S 3H8, Canada; hoekstra,gladders,hyee@astro.utoronto.ca\\\\ $^{3}$~Visiting Astronomer, Canada-France-Hawaii Telescope, which is operated by the National Research Council of Canada, Le Centre National de Recherche Scientifique, and the University of Hawaii\\\\ $^{4}$~Institut d'Astrophysique de Paris, 98 bis, Boulevard Arago, 75014, Paris, France\\\\ $^{5}$~Present address: Observatories of the Carnegie Institution of Washington, 813 Santa Barbara Street, Pasadena, California 91101\\\\ $^{6}$~Observatoire de Paris, LERMA, 61, Avenue de l'Observatoire, 75014, Paris, France }  These studies indicate that the bias parameter $b\\simeq 0.7$ on scales less than $\\sim 2 h_{50}^{-1}$ Mpc $(\\Omega_m=0.3,~\\Omega_\\Lambda=0.7)$. On larger scales $b$ increases to a value close to unity (Jenkins et al.  1998).  Although this procedure provides useful information about the bias parameter $b$ as a function of scale, it does rely on the assumptions made in the numerical simulations. In addition, it cannot be used to examine how tight the correlation between the matter and light distribution is. To do so, we need to measure the galaxy-mass cross-correlation coefficient $r$, which is a measure of the amount of stochastic and non-linear biasing (e.g., Pen 1998; Dekel \\& Lahav 1999; Somerville et al. 2001).  The coherent distortions in the shapes of distant galaxies caused by weak gravitational lensing provide a unique way to study the dark matter distribution in the universe. Weak lensing probes the dark matter distribution directly, regardless of the light distribution. In addition, it provides measurements on scales from the quasi-linear to the non-linear regime, where comparisons between observations and predictions are still limited. Much progress has been made in recent years, and the latest results give accurate joint constraints on the matter density $\\Omega_m$, and the normalisation of the power spectrum $\\sigma_8$ (e.g., Bacon et al. 2002; Hoekstra et al. 2002b; Refregier et al. 2002; van Waerbeke et al. 2002).  Although redshift surveys can be used to determine the relative values of $b$ and $r$ for different galaxy types (Tegmark \\& Bromley 1999; Blanton 2000), weak gravitational lensing provides the only direct way to measure the galaxy-mass cross-correlation function (e.g., Fischer et al. 2000; Wilson, Kaiser, \\& Luppino 2001; McKay et al. 2001; Hoekstra, Yee \\& Gladders 2001b).  Fischer et al. (2000) used the SDSS commissioning data to measure the galaxy-mass correlation function and their results suggested an average value of $b/r\\sim 1$ on submegaparsec scales. This approach has been explored by Guzik \\& Seljak (2001) who used semi-analytic models of galaxy formation combined with N-body simulations.  Their results suggest that the cross-correlation coefficient is close to unity. The galaxies used in their analysis, however, are massive (and consequently luminous) because of the limited mass resolution of the simulations.  In this Paper we use the method proposed by Schneider (1998) and Van Waerbeke (1998), which allows us to measure the biasing parameters directly as a function of scale. Hoekstra, Yee, \\& Gladders (2001a) applied this technique to 16 deg$^2$ of $R_C$-band imaging data from the Red-Sequence Cluster Survey (RCS; e.g., Yee \\& Gladders 2001), and measured the ratio $b/r$ as a function of scale. They found an average value of $b/r=1.05^{+0.12}_{-0.10}$ for a $\\Lambda$CDM cosmology on scales ranging from $150h_{50}^{-1}$~kpc to $3h_{50}^{-1}$~Mpc.  However, an accurate measurement of the mass auto-correlation function is required to constrain $b$ and $r$ separately. Furthermore the analysis is facilitated if the measurements of both the galaxy auto-correlation function and the mass auto-correlation function probe the power spectrum at the same redshift. To this end, the measurements presented by van Waerbeke et al. (2002), based on deep imaging data from the VIRMOS-DESCART survey, are ideal. In this Paper we combine the results of the RCS and the VIRMOS-DESCART survey, allowing us for the first time to measure both $b$ and $r$ as a function of scale.  The structure of the paper is as follows. In \\S 2 we discuss the data sets we use for our analysis. In \\S3 we define the bias parameters, and show how they can be related to observed correlation functions. The relevant correlation functions are discussed in \\S4. The actual measurements are discussed in \\S5.  The inferred bias parameters are presented in \\S6. In the appendix we demonstrate that the method used in this paper gives accurate results if the bias paramters vary with scale and redshift.  \\vspace{0.5cm} ",
        "conclusions": "The results presented in Figure~\\ref{ratios} suggest significant variation of both $b$ and $r$ with scale. We convert the observed ratios into estimates for the bias parameters for the currently favored cosmological model $(\\Omega_m=0.3,~\\Omega_\\Lambda=0.7)$, using Eqns.~14 and 17.  Figure~\\ref{bias}a shows the inferred value of the galaxy-mass cross-correlation coefficient $r$ as a function of scale. Because of the small value of $\\langle M_{\\rm ap}^2\\rangle$ at 1~arcminute, the uncertainties in both $b$ and $r$ are very large, and we have omitted this point.  We find that on small scales the measurements are consistent with a value $r\\sim1$. This good correlation between mass and light on small scales (i.e., around galaxies) indicates that (luminous) galaxies are surrounded by massive halos. We note that with our definition $r$ can be larger than 1. On the largest scales the results are consistent with $r=1$. On scales $\\sim 1h_{50}^{-1}$ Mpc $r$ is significantly lower than unity, with a minimum value of $r=0.57^{+0.08}_{-0.07}$ (68\\% confidence).  Figure~\\ref{bias}b shows that the scale dependence of $b$ is very similar to that of $r$ (also see Figure~\\ref{b_over_r}).  We find that $b$ is smaller than unity on scales $\\sim 1-2 h_{50}^{-1}$ Mpc, with a minimum value of $b=0.71^{+0.06}_{-0.05}$ (at $1.5 h_{50}^{-1}$ Mpc; 68\\% confidence). Hence, the dark matter is more strongly clustered than the galaxies. The variation of $b$ with scale is significant, but a better determination of $\\langle M_{\\rm ap}^2\\rangle$ is needed to study the scale dependence in more detail.  \\vbox{ \\vspace{-0.3cm} \\begin{center} \\leavevmode \\hbox{% \\epsfxsize=\\hsize \\epsffile[50 150 590 720]{f5.eps}} \\figcaption{\\footnotesize (a) The measured value of the galaxy-mass cross correlation coefficient $r$ as a function of scale for the $\\Lambda$CDM cosmology. (b) The bias parameter $b$ as a function of scale. The upper axis indicates the effective physical scale probed by the compensated filter $U(\\phi)$ at the median redshift of the lenses $(z=0.35)$. The errorbars correspond to the 68\\% confidence intervals. Note that the measurements at different scales are slightly correlated. \\label{bias}} \\end{center}}  Our result is in good qualitative agreement with the findings of Jenkins et al. (1998) who combined the observed APM two-point correlation function and the matter correlation function derived from numerical simulations.  We note that a direct comparison cannot be made with other measurements (such as those from Jenkins et al. 1998), because the bias properties depend on galaxy type and redshift. For instance, observations show that the amplitude of the galaxy correlation function depends on luminosity (e.g., Benoist et al. 1996; Norberg et al. 2001).  Our results are derived for a specific subset of galaxies, and are redshift averaged values (the lenses have a large range in redshift, with intrinsically brighter galaxies at higher redshifts). This is clearly demonstrated by Eqn~A7. In the simple case that the average bias parameters change approximately linearly with redshift, our results can be interpreted as a measurement of the biasing properties of $L_*$ galaxies at redshift $\\sim 0.35$. However, the redshift dependence of the bias in magnitude limited samples is usually more complex.  Hoekstra et al. (2001a) measured the ratio $b/r$ on angular scales out to 12.5 arcminutes, and found that the results were consistent with a constant value.  With the additional data (which probe larger scales and give smaller errorbars) we find evidence for a small trend with scale. The measurements of the current RCS data extend out to an angular scale of 1 degree (which corresponds to a physical scale of $12.5 h_{50}^{-1}$ Mpc), and are presented in Figure~\\ref{b_over_r}a. Figure~\\ref{b_over_r} shows that even on a 1 degree scale, the systematics are much smaller than the signal. Hence, the value of $b/r$ can be determined accurately. The new data yield an average value of $b/r=1.090\\pm0.035$. Limiting the measurements to the angular scales studied in Hoekstra et al. (2001a), we find an average value of $b/r=1.05\\pm0.04$, in excellent agreement with the previous estimate of $b/r=1.05^{+0.12}_{-0.10}$.   Our measurements should be compared to models of galaxy formation. The two commonly used approaches are hydrodynamic simulations (e.g., Blanton et al. 2000, Yoshikawa et al. 2001) or semi-analytical models (e.g., Kauffmann et al. 1999a,b; Somerville et al. 2001, Guzik \\& Seljak 2001). These studies suggest values of $r$ close to unity. Unfortunately, the mass resolution in these simulations, such as the GIF results (e.g., Kauffmann et al. 1999a, b), is too poor: they only resolve massive, luminous galaxies. Our sample of lenses contains many lower luminosity systems, which seriously hampers the comparison of our results with predictions.  As the value of $r$ is intimately linked to the details of galaxy formation, a careful comparison of the models with the weak lensing measurements provides unique constraints. In addition, planned large weak lensing surveys, such as the CFHT Legacy Survey\\footnote{\\tt http://www.cfht.hawaii.edu/Science/CFHLS/}, will significantly improve the accuracy of the measurements, allowing larger scales to be probed, and enabling us to select galaxies based on the colours or luminosities (using photometric redshifts). Therefore the prospects of constraining biasing parameters from weak lensing are excellent.   \\vbox{ \\vspace{-0.3cm} \\begin{center} \\leavevmode \\hbox{% \\epsfxsize=8cm \\epsffile[50 150 590 720]{f6.eps}} \\figcaption{\\footnotesize (a) The value of $b/r$ as a function of angular scale, under the assumption $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$. These results are based on RCS data only, which allows this ratio to measured out to 1 degree (which corresponds to a physical scale of $12.5 h_{50}^{-1}$ Mpc). Note that the points are slightly correlated. The value of $b/r$ is almost constant over the range probed here. For this cosmology we find an average value of $b/r=1.090\\pm 0.035$, in excellent agreement with the result from Hoekstra et al. (2001a). (b) Measurement of $\\langle M_{\\rm X}{\\cal N}\\rangle/\\langle{\\cal N}^2\\rangle$ (the signal when the phase of the shear is increased by $\\pi/2$), which should vanish if the results in (a) are caused by lensing. \\label{b_over_r}} \\end{center}}"
    },
    "0206/astro-ph0206429_arXiv.txt": {
        "abstract": "{ We report on a detailed study of the X-ray spectrum of the nearby eclipsing spectroscopic binary \\yygem. Observations were obtained simultaneously with both large X-ray observatories, {\\em XMM-Newton} and {\\em Chandra}. We compare the high-resolution spectra acquired with the Reflection Grating Spectrometer onboard {\\em XMM-Newton} and with the Low Energy Transmission Grating Spectrometer onboard {\\em Chandra}, and evidence in direct comparison the good performance of both instruments in terms of wavelength and flux calibration. The strongest lines in the X-ray spectrum of \\yygem are from oxygen. Oxygen line ratios indicate the presence of a low-temperature component ($1-4$\\,MK) with density $n_{\\rm e} \\leq 2\\,10^{10}\\,{\\rm cm^{-3}}$. The X-ray lightcurve reveals two flares and a dip corresponding to the secondary eclipse. An increase of the density during phases of high activity is suggested from time-resolved spectroscopy. Time-resolved global fitting of the European Photon Imaging Camera CCD spectrum traces the evolution of temperature and emission measure during the flares. These medium-resolution spectra show that temperatures $> 10^7$\\,K are relevant in the corona of YY\\,Gem although not as dominant as the lower temperatures represented by the strongest lines in the high-resolution spectrum. Magnetic loops with length on the order of $10^9$\\,cm, i.e., about 5\\,\\% of the radius of each star, are inferred from a comparison with a one-dimensional hydrodynamic model. This suggests that the flares did not erupt in the (presumably more extended) inter-binary magnetosphere but are related to one of the components of the binary. ",
        "introduction": "\\label{sect:intro}  The Castor system comprises three visual stars, all of which are spectroscopic binaries. The optically faintest of the three components, \\yygem (= Castor C), is of major importance to stellar evolution studies: being an eclipsing spectroscopic binary ($i \\sim 86^\\circ$; \\cite{Pettersen76.1}) very near or on the main sequence it allows for the determination of masses and radii of both components, and for a test of evolutionary tracks. Both stars in the \\yygem system are of nearly equal spectral type, dM1e, in a synchronous orbit with a period of 0.81\\,d. \\yygem was the first late-type star on which periodic photometric variability was detected (\\cite{Kron52.1}). \\citey{Bopp73.1} have explained similar brightness variations on the flare star BY\\,Dra by rotation of star spots, and stars displaying this phenomenon are since then termed BY\\,Dra variables. Doppler images of \\yygem have revealed spots at mid-latitudes on both stars (\\cite{Hatzes95.1}).  X-rays from the Castor sextuplet were first recorded by the {\\em Einstein} satellite (\\cite{Vaiana81.1}, \\cite{Caillault82.1}, \\cite{Golub83.1}). Since then the system was observed by virtually all X-ray observatories; see e.g., \\citey{Pallavicini90.1}, \\citey{Gotthelf94.1}, \\citey{Schmitt94.1}, \\citey{Guedel01.1}. As the spatial resolution of the instruments improved, X-ray emission from individual members of the Castor system could be identified. At present, all three visual binaries in the sextuplet are known to be X-ray sources. The X-ray emission from Castor\\,A and~B is commonly attributed to their late-type companions because the primaries, being early A-type stars without a convective envelope and consequently with no ability to support the dynamo mechanism, are not expected to be X-ray sources.  In nearly all of the previous X-ray studies of \\yygem flare activity was reported. The luminosity of \\yygem in the soft X-ray band ranges between $\\sim (2 - 8)\\,10^{29}\\,{\\rm erg/s}$. Flares have been observed also in other parts of its electromagnetic spectrum. \\citey{Doyle90.1} found that the flare activity on \\yygem is more than an order of magnitude larger during out-of-eclipse times as compared to times when one of the stars is eclipsed. On the basis of these U band observations it was suggested that the amount of magnetic energy is largest in the inter-binary space, leading to frequent energy release in this region. A periodicity of the optical flaring rate was reported by \\citey{Doyle90.2}, but has not been confirmed so far. The chromosphere and transition region was examined by \\citey{Haisch90.1}, and moderate variability throughout the orbital cycle was detected in the emission lines.  The extraordinary variability of \\yygem has instigated us to carry out a coordinated multi-wavelength campaign comprising radio, optical and X-ray observations. As a unique example of an eclipsing spectroscopic binary with nearly identical components \\yygem is ideally suited for a study of stellar coronal structure. In particular we were aiming at looking into the possibility of enhanced flare activity in the inter-binary space as suggested by \\citey{Doyle90.1}. Magnetic structures in between the components of binaries have been discussed also for RS\\,CVn systems (e.g., \\cite{Uchida83.1}). Additional information about the coronal geometry can be obtained from observations of X-ray eclipses. The depth and duration of the eclipses provide a measure for the spatial extent of the corona and contribution of the two components in the binary. From a 3D-deconvolution of the X-ray lightcurve of \\yygem (using an earlier observation by {\\em XMM-Newton}) \\citey{Guedel01.1} found that the coronae of both stars in the binary are inhomogeneously structured with brighter areas at mid-latitudes. Obscuration of coronal features can also be used to localize and constrain the emitting region in active stars (see e.g., Gunn et~al. 1997, 1999, \\cite{Schmitt99.1}).  As part of our multi-wavelength project the Castor system was observed simultaneously with both {\\em XMM-Newton} and {\\em Chandra}. While in the EPIC MOS image presented by \\citey{Guedel01.1} Castor~A and~B were spatially resolved, our {\\em Chandra} observation allows for the first time to separate the X-ray spectra of Castor~A and~B from each other. We will examine the X-ray characteristics of Castor~A and~B in a related publication (Stelzer et al., in prep.). Here we concentrate on the intermediate and high-resolution X-ray spectrum of YY\\,Gem. We emphasize that this observation presents an excellent opportunity to cross-check the calibration of the dispersive instruments onboard both satellites. We analyse the emission line spectrum observed with the {\\em Chandra} Low-Energy Transmission Grating Spectrometer (LETGS) and the Reflection Grating Spectrometer (RGS) onboard {\\em XMM-Newton} independently of each other, and derive observed line fluxes for both instruments. Emission lines in the spectral region of the LETGS and RGS include the helium-like triplets from C\\,V to Si\\,XIII, the Lyman series of hydrogen-like ions, and numerous iron L-shell transitions. By means of time-resolved analysis of spectral parameters we check for spectral variability related to changes in the activity level of YY\\,Gem. We use both `global' fits to the EPIC X-ray spectrum from $0.3$ to $10$\\,keV and modeling of individual emission lines in the range 2-175\\,\\AA~~ (resolved by the gratings onboard both {\\em Chandra} and {\\em XMM-Newton}) to infer the time evolution of temperature and density of the emitting plasma.  The {\\em XMM-Newton} and {\\em Chandra} observations are introduced in Sect.~\\ref{sect:obs}. In Sect.~\\ref{sect:highres-spectra} we present results from high-resolution spectroscopy, involving a detailed comparison of the contemporaneous LETGS and RGS spectrum. Time-resolved spectroscopy is discussed in Sect.~\\ref{sect:timeres}. We draw conclusions on the coronal structure by means of loop modeling (Sect.~\\ref{sect:loops}), and summarize the results in Sect.~\\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}  In September 2000 \\yygem was observed simultaneously by {\\em XMM-Newton} and {\\em Chandra} providing a possibility to cross-check the performance of the instruments on both satellites. We found that the line centers and the line fluxes measured by LETGS and RGS are in good agreement with only few exceptions. One of these is the underestimation of the O\\,VII resonance line flux due to an absorption edge in the RGS which was not yet taken account of in the response matrix at the time of analysis (\\cite{denHerder02.1}): For the resonance line of the O\\,VII triplet the RGS flux is found to be $\\sim 70$\\% of the LETGS flux, but the measurements of the two instruments are compatible with each other at the 3\\,$\\sigma$ level. We note that similar differences have been observed between LETGS and RGS for Procyon (see \\cite{Raassen02.1}). A deviation of the flux for the Ne\\,X Ly$\\alpha$ line between the two sides of the LETGS seems to point at a problem on the right side of this detector at the respective wavelength (12.13\\,\\AA).  The high-resolution spectrum of \\yygem is dominated by oxygen and neon lines and lines from intermediate ionization stages of iron, generally Fe\\,XVIII and below. Only few lines from highly ionized iron (Fe\\,XIX and higher) are seen at the long wavelength end of the LETGS spectrum. The characteristic Fe\\,XXI line at 128.73\\,\\AA~ is detected but very weak, underlining that the corona of \\yygem~ is not dominated by the high temperatures ($\\geq 4\\,$MK) typically observed on RS\\,CVn binaries (see e.g., \\cite{Mewe01.1}, \\cite{Audard01.1}).  We have probed the optical depth in the corona of YY\\,Gem using characteristic ratios of Fe\\,XVII line fluxes. Comparing these ratios with measurements for other stars known to exhibit very different activity levels shows that optical depth effects do not play a role in stellar coronae of all types.  The correlation between X-ray luminosity and temperature in stellar coronae is well known (see e.g., \\cite{Guedel97.1}). Now, the dispersive instruments of {\\em Chandra} and {\\em XMM-Newton} allow to confirm this relation using intensity ratios of individual emission lines. % \\citey{Ness02.2} show in their comparison of {\\em Chandra} observations of ten stars that \\yygem fits well into the relation between $L_{\\rm x}$ and temperature probed by the O\\,VIII to O\\,VII flux ratio representing a moderately active star.  The only element in the spectrum of YY\\,Gem with a sufficiently strong and unblended helium-like triplet is oxygen. The $R = f/i$ value derived from the observation presented here is compatible with an earlier measurement by \\citey{Guedel01.1}. For the time-averaged spectrum the ratios of the oxygen triplet yield a temperature of $T = 2-3$\\,MK and an upper limit to the electron density of $n_{\\rm e} \\leq 2\\,10^{10}\\,{\\rm cm^{-3}}$. During the time of observation \\yygem showed strong variability including two flares. % The time-resolved analysis of the high-resolution spectrum reveals variations of $R$ which are suggestive of an increase in density during times of enhanced count rate. Our measurement is one of the first direct indications for the increase of the coronal density during flares on stars other than the Sun derived from high-resolution X-ray spectroscopy. A density measurement in the corona of AB\\,Dor using an {\\em XMM-Newton} measurement of oxygen triplet lines, e.g., resulted in comparable densities during flare and during quiescence (\\cite{Guedel01.2}). A detailed investigation of line ratios in the short time intervals representative for the different activity levels of \\yygem is constricted by low S/N. The time-resolved high-resolution spectroscopy, furthermore, shows that during flares the absolute amount of (oxygen emission by) relatively low-temperature plasma increases. Lines from elements probing higher temperatures more typical for flares were not strong enough for this kind of analysis.  The presence of higher temperatures in the corona of \\yygem is, however, evident from our time-resolved spectroscopy with EPIC. The temperatures for the global fit of the quiescent EPIC spectrum (2.4\\,MK, 7.4\\,MK, and 20.8\\,MK) comprise the temperature range measured from the analysis of oxygen lines, but go well beyond these values. Two additional VMEKAL components are necessary to describe the EPIC spectrum during flares. The corresponding temperatures and emission measures are highest during the rise phase. An attempt to reproduce the EPIC spectrum with a number of Gaussian profiles using the information from the RGS spectrum was described by \\citey{Haberl02.1}. But the low resolution of EPIC makes this instrument insensitive to changes in the relative strength of nearby emission lines during different phases of activity.  With plasma parameters derived from spectral modeling % the size of the (flare) emitting structures can be inferred. We used a hydrodynamic approach (see \\cite{Reale97.1}) involving heating during the decay phase  % to determine the length of flaring loops on YY\\,Gem. The overall size of quiescent stellar coronae can be estimated making use of plasma density and temperature derived from the analysis of high-resolution spectra (see e.g., \\cite{Ness01.1}, \\cite{Mewe01.1}, and \\cite{Raassen02.1}). However, loop lengths derived from individual lines refer only to a very limited temperature range and the relevance of the single loop model for the quiescent corona is questionable. In our case the spectral parameters derived from individual lines are not useful for this kind of analysis, since the upper limits on the density translate to lower limits of the loop size according to the RTV-laws. Instead, we derived loop sizes using results from global fitting with EPIC. We find that loops are small ($2\\,10^9$\\,cm), on the order of a few percent of the stellar radius ($R_* \\sim 0.6\\,R_\\odot$) of each of the components in the YY\\,Gem binary. This is consistent with the compactness of the corona established from eclipse mapping for the earlier {\\em XMM-Newton} observation of YY\\,Gem (\\cite{Guedel01.1}). Small loops point at magnetic activity on the individual stars rather than large scale structures associated with overlapping magnetospheres. Furthermore, the RTV-laws point at high densities in the loop, confirming the evidences for low $R-$ratio during times of flares."
    },
    "0206/gr-qc0206002_arXiv.txt": {
        "abstract": "We study the dynamical instability against bar-mode deformation of differentially rotating stars. We performed numerical simulation and linear perturbation analysis adopting polytropic equations of state with the polytropic index $n=1$. It is found that rotating stars of a high degree of differential rotation are dynamically unstable even for the ratio of the kinetic energy to the gravitational potential energy of $O(0.01)$. Gravitational waves from the final nonaxisymmetric quasistationary states are calculated in the quadrupole formula. For rotating stars of mass $1.4M_{\\odot}$ and radius several 10 km, gravitational waves have frequency several $100$ Hz and effective amplitude $\\sim 5 \\times 10^{-22}$ at a distance of $\\sim 100$ Mpc. ",
        "introduction": "Stars in nature are rotating and subject to nonaxisymmetric rotational instabilities.  An exact treatment of these instabilities exists only for incompressible equilibrium fluids in Newtonian gravity \\cite{CH69,TA78}. For these configurations, global rotational instabilities arise from nonradial toroidal modes $e^{im\\varphi}$ ($m=\\pm 1,\\pm 2, \\dots$) when $\\beta\\equiv T/W$ exceeds a certain critical value. Here $\\varphi$ is the azimuthal coordinate and $T$ and $W$ are the rotational kinetic and gravitational potential energies.  In the following we will focus on the $m=\\pm 2$ bar-mode since it is expected to be the fastest growing mode (but see Centrella et al. (2001) with regard to a counter-example for extremely soft equations of state).  There exist two different mechanisms and corresponding timescales for bar-mode instabilities.  Uniformly rotating, incompressible stars in Newtonian theory are {\\em secularly} unstable to bar-mode deformation when $\\beta \\geq \\beta_s \\simeq 0.14$.  However, this instability can only grow in the presence of some dissipative mechanism, like viscosity or gravitational radiation, and the growth time is determined by the dissipative timescale, which is usually much longer than the dynamical timescale of the system.  By contrast, a {\\em dynamical} instability to bar-mode deformation sets in when $\\beta \\geq \\beta_d \\simeq 0.27$.  This instability is independent of any dissipative mechanisms, and the growth time is determined typically by the hydrodynamical timescale of the system.  For the compressible case, determining the onset of the dynamical bar-mode instability, as well as the subsequent evolution of an unstable star, requires numerical computations. Hydrodynamical simulations performed in Newtonian theory \\cite{TDM,DGTB,WT,TH,CT2,CT3,CT1,PDD,Toman1,Toman2,New,brown,LL1,LL2} have shown that $\\beta_d$ is $\\sim 0.27$ as long as the rotational profile is not strongly differential. In this case, once a bar has developed, the formation of spiral arms plays an important role in redistributing the angular momentum and forming a core-halo structure. Recently, it has been shown that $\\beta_d$ can be smaller for stars with a higher degree of differential rotation \\cite{CNLB,TH,PDD,LL1,LL2} as $\\beta_d \\sim 0.14$. In such case, the formation of bars and spiral arms does not take place. Instead, small density perturbation is excited and left to be weakly nonlinear. To date, there is no report of dynamically unstable stars with $\\beta \\alt 0.14$.  There are numerous evolutionary paths which may lead to the formation of rapidly and differentially rotating neutron stars. $\\beta$ increases approximately as $R^{-1}$ during stellar collapse. Also, with decreasing stellar radius, the angular velocity in the central region may be much larger than that in the outer region. During supernova collapse, the core contracts from $\\sim 1,000$ km to several 10 km, and hence $\\beta$ increases by two orders of magnitude. Thus, even rigidly rotating progenitor stars with a moderate angular velocity may yield rapidly and differentially rotating neutron stars which may reach the onset of dynamical instability \\cite{BM,RMR}. Similar arguments hold for accretion induced collapse of white dwarfs to neutron stars and for the merger of binary white dwarfs to neutron stars. Differential rotation is eventually suppressed by viscous angular momentum transfer or magnetic braking \\cite{BSS}, but within the transport timescale, the differentially rotating stars are subject to nonaxisymmetric instabilities.  In this Letter, we report some of results we have recently obtained on dynamical bar-mode instabilities in differentially rotating stars. We pay attention to rotating stars with a high degree of differential rotation in Newtonian gravity using both linear perturbation analysis and nonlinear hydrodynamical simulation. The detail of our study will be summarized in a future paper \\cite{SKE}. Here, we highlight a new finding. ",
        "conclusions": "Using the results shown in Fig. 6, we can estimate the expected amplitude of gravitational waves from the nonaxisymmetric outcome formed after the dynamical instability saturates. Here, we pay particular attention to protoneutron stars likely formed soon after the supernovae of mass $\\sim 1.4M_{\\odot}$ and radius several 10 km. For the maximum amplitude of gravitational waves $h \\sim 0.2M^2/Rr$, and the energy luminosity $\\dot E \\sim 0.01(M/R)^5$, the emission timescale of gravitational waves can be estimated as \\beq \\tau \\sim {T \\over \\dot E} \\sim 100 \\beta' \\biggl({R \\over M}\\biggr)^4 M, \\eeq where we set $T = \\beta' M^2/R$ and $0.04 \\alt \\beta' \\alt 0.4$ for $\\hat A=0.3$ and $0.3 \\leq C_a \\leq 0.8$. The characteristic frequency of gravitational waves is denoted as \\beqn f \\equiv {\\omega_r \\over 2\\pi} \\approx 800 {\\rm Hz}~ \\biggl({\\xi \\over 2}\\biggr) \\biggl({15M \\over R}\\biggr)^{3/2} \\biggl({M \\over 1.4M_{\\odot}}\\biggr)^{-1}, \\eeqn where $\\xi\\equiv\\omega_r\\sqrt{R^3/M} \\sim 2$ (see Figs. 4 and 7). Since the cycles of gravitational wave-train $N$ are estimated as $N \\equiv f\\tau$, the effective amplitude of gravitational waves $h_{\\rm eff} = \\sqrt{N}h $ is % \\beqn h_{\\rm eff} \\approx 5 \\times 10^{-22} \\biggl({\\xi \\over 2}\\biggr)^{1/2} \\biggl({\\beta' \\over 0.1}\\biggr)^{1/2} \\biggl({R \\over 15M}\\biggr)^{1/4} \\nonumber \\\\ \\hskip 2cm \\times \\biggl({M \\over 1.4M_{\\odot}}\\biggr) \\biggl({100{\\rm Mpc} \\over r}\\biggr) % \\eeqn \\cite{Kip,LS,LL1,LL2}. Thus, gravitational waves from protoneutron stars of a high degree of differential rotation and of mass $\\sim 1.4 M_{\\odot}$ and radius several $10$ km in the distance of $\\sim 100$ Mpc can be a source for laser interferometric detectors such as LIGO \\cite{Kip2}. We emphasize that $f$ and $h_{\\rm eff}$ depend weakly on $\\beta$ (or $\\beta'$) for $\\hat A=0.3$. This implies that even if the star is not rapidly rotating, the dynamical instability can set in and as a result, differentially rotating stars can emit gravitational waves of a large amplitude.    \\begin{figure}% \\vspace*{-5mm} \\begin{center} \\leavevmode \\psfig{file=fig7.ps,width=2.5in,angle=0} \\end{center} \\vspace*{-20mm} \\caption{$\\Omega_0/\\sqrt{M/R^3}$ as a function of $\\beta$. The solid circles denote the dynamically unstable stars we found. } \\end{figure}    To summarize, we have studied dynamical bar-mode instability of differentially rotating stars focusing on the f-mode. We have found that rotating stars of a high degree of differential rotation are dynamically unstable against nonaxisymmetric deformation even for $\\beta \\ll 0.27$.   It is worthy to note that the real parts of the eigen frequencies do not vanish but approach to finite values as the value of $\\beta$ decreases, i.e. in the spherical limit. This behavior is totally different from those of the r-modes (see e.g. Karino et al. 2001), several self-gravity induced instability modes of slender tori or annuli such as I-modes and J-modes \\cite{GN,AT}, or shear instability modes (or P-modes) such as Papaloizou-Pringle instability for toroids \\cite{PP} and for spheroids \\cite{Lu}. Therefore, we have identified the unstable modes we find in this paper as the f-mode.  The physical mechanism for the onset of dynamical instabilities found in this Letter may be explained in the following manner: For a small value of $\\hat A$, non-spherical deformation of a stellar configuration is significant around the rotational axis although the distant part from the rotational axis is almost spherical. Thus, with decrease of the value of $C_a$, most rotational energies are confined to the region near the rotational axis. Increase of the rotational energy can be much larger than those of the gravitational and internal energies, because the overall shape cannot be much different from that of a sphere. As a result, the total energy becomes large with small values of $C_a$. However, if the rotational energy exceeds a certain amount in the region near the rotational axis, there may exist other equilibrium configurations with lower total energies, as in the case of bifurcation of the Jacobi ellipsoidal sequence from the Maclaurin sequence \\cite{CH69}. In fact, the total energy of the Jacobi ellipsoid is lower than the Maclaurin spheroid of the same mass and the angular momentum, if the rotational energy exceeds a certain criterion.  Once the instability sets in, the axisymmetric star begins to change its shape to a nonaxisymmetric configuration. In such case, the angular velocity in the region near the rotational axis, $\\sim \\Omega_0$, is very large (larger than $\\sqrt{M/R^3}$). It implies that, despite of the large rotational energy, the angular momentum cannot be very large because $ T \\sim \\int (x^2 + y^2) \\times \\Omega_0^2 dm $ and $J \\sim \\int (x^2 + y^2) \\times \\Omega_0 dm$, where $J$ and $dm$ are the angular momentum and the mass element, respectively. Therefore, even if the nonaxisymmetric mode grows due to the dynamical instability, the configuration cannot be highly elongated or form spiral-arm structures; i.e., there exists a saturation state of a small nonaxisymmetric deformation. For the bar-mode instability of the Maclaurin spheroids, the rotational energy is \"confined\" to the outer part of the configuration by making the shape very flat and so the angular momentum can be also large, even though the angular velocity is not extremely large ($<\\sqrt{M/R^3}$). Therefore, once a bar-mode dynamical instability sets in, configurations can be considerably different from the original spheroidal shapes.     It is possible that other non-fundamental modes have larger $\\omega_i$ than the f-mode has. However, the present study shows that there exists at least one mode which induces the nonaxisymmetric deformation. Such rotating stars subsequently form nonaxisymmetric structures and hence can be a source of laser interferometric gravitational wave detectors."
    },
    "0206/astro-ph0206209_arXiv.txt": {
        "abstract": "{We present a model for understanding the origins and evolution of galaxy morphologies from a phenomenological perspective.  The model includes an observationally motivated prescription for star formation in galaxy disks, as well as a merger-driven starburst mode of star formation.  We consider the formation of bulges and ellipticals both in mergers and by global instabilities in disks.  We use our model to investigate the fundamental properties of disk galaxies and the effects of surface-brightness limits on large galaxy surveys. }   \\addkeyword{Galaxies: formation}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  Our understanding of galaxy formation in a cosmological context has increased considerably over the last decade through the use of phenomenological models.  In these models, stars form primarily in galaxy disks and bulges are assumed to form when two or more disk galaxies merge, or when bar instabilities lead to vertical heating of the disk.  We use such a model for the evolution of galaxy morphologies to investigate the statistics of the fundamental parameters of disk galaxies.  These results are particularly important for the interpretation of magnitude- or diameter-limited galaxy surveys, which may miss a significant population of low-surface-brightness (LSB) galaxies. ",
        "conclusions": ""
    },
    "0206/astro-ph0206386_arXiv.txt": {
        "abstract": "High-quality spectropolarimetric data (range 417-860 nm; spectral resolution 1.27 nm and 0.265 nm/pixel) of SN 2002ap were obtained with the ESO Very Large Telescope Melipal (+ FORS1) at 3 epochs that correspond to -6, -2, and +1 days for a V maximum of 9 Feb 2002. The polarization spectra show three distinct broad ($\\sim$ 100 nm) features at $\\sim$ 400, 550, and 750 nm that evolve in shape, amplitude and orientation in the Q-U plane. The continuum polarization grows from nearly zero to $\\sim$ 0.2 percent. The 750 nm feature is polarized at a level $\\gta$ 1 percent. We identify the 550 and 750 nm features as Na I D and OI $\\lambda$ 777.4 moving at about 20,000 \\kms. The blue feature may be Fe II.  We interpret the polarization evolution in terms of the impact of a bipolar flow from the core that is stopped within the outer envelope of a carbon/oxygen core.  Although the symmetry axis remains fixed, as the photosphere retreats by different amounts in different directions due to the asymmetric velocity flow and density distribution geometrical blocking effects in deeper, Ca-rich layers can lead to a different dominant axis in the Q-U plane.  We conclude that the features that characterize SN~2002ap, specifically its high velocity, can be accounted for in an asymmetric model with a larger ejecta mass than SN~1994I such that the photosphere remains longer in higher velocity material.  The characteristics of ``hypernovae\" may be the result of orientation effects in a mildly inhomogeneous set of progenitors, rather than requiring an excessive total energy or luminosity.  In the analysis of asymmetric events with spherically symmetric models, it is probably advisable to refer to ``isotropic equivalent\" energy, luminosity, ejected mass, and nickel mass. ",
        "introduction": "From the first identification of Type~Ic supernovae (SN~Ic) as a distinguishable spectral subclass \\citep{WheelLev85} it has been recognized that they represent the collapse of bare non-degenerate, carbon/oxygen cores \\citep{vanDyk:1992,CloccWheel97}.  Most display a range in peak absolute magnitudes from -16.5 to -18.5, but some may be considerably brighter, exceeding -19 \\citep{Clocc:2000} without showing excessive velocities or light curve anomalies.  They also show a variety of rates of decay from maximum \\citep{Clocc83N:1996, Clocc83V:1997, CloccWheel97, CloccIcLC97}. The best guess is that they represent the collapse of the core of a massive star that has lost most or all of both its hydrogen and helium layers.  The remaining core mass and hence ejecta mass is expected to show some dispersion that will be reflected in the properties of the observed explosions. In addition, the suspicion that all core collapse supernovae, and even more so all SN~Ic, are strongly asymmetric \\citep{Wang:2001} suggests that there will be significant observer line-of-sight effects that can be manifest in the luminosity and photospheric velocity as well as the spectropolarimetry.  The observed properties, especially the photospheric velocity, are also expected to be strong functions of epoch with higher velocities at earlier times.  The advent of SN~1998bw and its apparent connection with GRB~980425 \\citep{Galama98} drew new attention to the general class of SN~Ic. SN~1998bw displayed large photospheric velocities early on and was a strong source of radio radiation implying some relativistic ejecta \\citep{Kulkarni98}, although an especially weak \\grb. At about the same time \\citet{Pacz98} coined the term ``hypernova\" to specifically mean the events with large optical luminosity that were associated with \\grbs\\ and their afterglows.  This term was adopted by some in the supernova community to represent supernovae, SN~1998bw in particular, that seemed to be excessively bright and to require very large amounts of kinetic energy and ejected \\ni\\ mass \\citep{Iwamoto98bw, Woosley98bw}.  Since that time, the term ``hypernova\" has been used to apply to events that resemble SN~Ic, but have large photospheric velocities, whether or not they are especially bright \\citep{Iwamoto:97ef, Mazef00}.  These developments raise the issue of whether ``hypernovae\" as the term is applied to hydrogen-deficient supernovae is really a separate class with excessively large kinetic energy and ejected \\ni\\ mass or whether these supernovae can be explained in terms of the normal range of properties, including asymmetries, associated with bare carbon/oxygen core collapse \\citep{HWW:1999}.  SN~2002ap has attracted much attention because early spectra showed a lack of hydrogen and helium characteristic of SN~Ic and broad velocity components \\citep{KinugasaIAU, MeikleIAU,Gal-YamIAU}, which, as mentioned above, have been taken as one characteristic of ``hypernovae.\" The nature, existence of, and import of ``hypernovae\" remains to be clarified, and the study of SN~2002ap presents an important opportunity. This supernovae was not associated with a \\grb, was not especially bright \\citep{Gal-Yamap02, Mazap02}, and was a weak radio source \\citep{Berger:2002}. Since its only special property seems to be a high photospheric velocity at early times, study of the velocity in the context of geometrical asymmetries may shed light on the general category of ``hypernovae.\"  A particularly important issue is to understand whether asymmetries could be affecting the interpretation of the kinematics, luminosity, kinetic energy, nickel mass, ejected mass, and progenitor mass of the supernova.  The prime tool to investigate the geometry of distant supernovae is spectropolarimetry \\citep{Wang:1996,Wang:2001, Leonard:2000,Leonard:2001a,Leonard:2001b,Howell:2001}.  Here we present spectropolarimetry of SN~2002ap from about six days before to about one day after peak visual magnitude.  This is the first time that the temporal evolution of the polarization has been studied in a Type Ib/c and the first time that such early polarimetry has been obtained of a Type Ib/c.  \\citet{Wang:1996} studied supernova polarimetry published before 1996 and found that Type Ia supernovae are normally not polarized and on average show much lower polarization than core-collapse supernovae (Type II, Ib/c; but see \\citealt{Howell:2001}). Subsequent data indicated a higher degree of polarization for the bare-core or low mass events such as Type IIb, and Ib/c supernovae, than for more massive ejecta \\citep{Wang:2001}. The degree of polarization is a function of time after explosion. For core-collapse events, the degree of polarization is higher past optical maximum than before optical maximum \\citep{Wang:2001,Leonard:2001a,Leonard:2001b}. \\citet{Wang:2001} reported the polarization of the Type Ic SN~1997X that was exceptionally high, around 7 percent. Some of this is probably interstellar in origin, but the suggestion remains that ``bare core\" supernovae like SN Ic are very highly distorted.  There was only one epoch of data on SN~1997X.  SN~2002ap gives the opportunity to study the polarization evolution of a Type Ib/c event in unprecedented detail.  We present in \\S 2 observations and data reduction of SN 2002ap.  We discuss the spectroscopic and spectropolarimetric evolution in \\S 3. The structure of the SN 2002ap ejecta is studied in \\S 4. In \\S 5, we give some discussion and conclusions with emphasis on implications for models of SN Ib/c and ``hypernovae.\" ",
        "conclusions": "Understanding the 3-D nature of core collapse supernovae, especially ``naked core\" events is an important key to the underlying physics, and polarimetry is the best tool to probe multi-dimensional effects.  SN~2002ap resembles other Type Ib/c events, but displays rather broad lines near maximum.  The distinct polarization spectra show that the ejecta are asymmetric.  The premaximum polarization spectra reveal a strong feature of O I 777.4 nm and probably Na I D, both moving at about 20,000 \\kms.  Near maximum light, a new dominant axis is revealed as defined, most likely, by calcium again with a velocity of about 20,000 \\kms.   These observations can be accounted for in a model in which a bi-polar flow from the exploding core imposes a predominantly prolate velocity structure and oblate density structure.  The outer photosphere at early times can thus have a complex structure with an effective dominant axis that is intermediate between the axis of symmetry and the equator.  Later observations would reveal a more predominantly oblate structure oriented along the equatorial plane, as observed.  This picture is supported by the lack of observations in the NIR of freshly synthesized nickel and cobalt as seen in Type Ia, suggesting that the bi-polar flow has induced velocity and density asymmetries, but has not penetrated the photosphere at early times so that these freshly synthesized elements are not observed.  We find no evidence for substantial amounts of matter moving at 45,000 \\kms, which was mentioned as a possibility by \\citet{Wang:ap02IAU}. The relatively high velocities that are observed, about 20,000 \\kms, broaden the lines compared to some other Type Ic, for instance SN~1994I \\citep{Iwamoto:94I}.  This can be accounted for by a combination of different ejecta mass and an asymmetric ejecta. In general, a higher ejecta mass and perhaps a steeper density gradient at the photosphere, can prolong the phase in which the photosphere resides in higher velocity material. We note also that the earliest spectra of SN~2002ap were obtained substantially before maximum light, a rare happenstance for SN~Ic events.  In addition to this basic effect of the ejecta mass, the asymmetry of the ejecta will bring line-of-sight effects.  As shown by \\citet{HWW:1999} the isodensity contours can be rather irregular, with a quadrupole-like distortion as the ejecta expand more rapidly along the poles, even for an ellipsoidal density distribution.  For a more realistic bi-polar explosion the density contours can be even more complicated.  Thus there may be special viewing angles along which the homologous expansion is relatively slow, leading to a delayed recession of the photosphere and hence to a larger line-of-sight photospheric velocity at a given epoch.  In any case, it is clear that while a proper interpretation of polarized, asymmetric ejecta is tricky, it is certainly risky to interpret observations of asymmetric ejecta in terms of a spherically symmetric model and to take the results literally.  The luminosity of SN~2002ap was consistent with a normal Type Ib/c, there was no \\grb, and the radio observations are consistent with only a modest ejection of high-velocity matter \\citep{Berger:2002}. The relatively high early photospheric velocities of SN~2002ap with respect to other SN~Ic, may then just be a question of epoch of observation, mass of ejecta, and orientation angle of the observer, and not a hint of exceptional properties.  If one can interpret SN~2002ap as a variation on the ``normal\" theme of Type Ib/c supernovae, then perhaps one should reconsider the general category of ``hypernovae.\"  In the context of supernovae (as opposed to Paczy\\'nski's original definition), SN~19998bw is the prototype of this proposed new category with large kinetic energy, ejecta mass, nickel mass and luminosity attributed to it in the context of spherically-symmetric models.  \\citet{Iwamoto98bw} determined that the exploding carbon/oxygen core had a mass of 12 to 15 \\Msun\\ with kinetic energy in the range 20 to 50$\\times10^{51}$ erg and mass of \\ni\\ in the range 0.6 to 0.8 \\Msun. \\citet{Woosley98bw} favored a carbon/oxygen core of 6 \\Msun, a kinetic energy of $22\\times10^{51}$ erg and a mass of \\ni\\ of 0.5 \\Msun.  Both of these models had some problems reproducing the spectra or light curves. This class of models also had problems reproducing the late time tails.  The predicted light curves tended to drop too quickly.  Another intrinsic problem is that these spherical models with very high kinetic energy imposed at the base of the ejecta predict very high minimum velocities.  The latter two problems are related.  The rapid expansion of the innermost layers cause them to thin out, thus decreasing the \\gr\\ deposition and causing the drop in the light curve.  By contrast, a somewhat schematic model by \\citet{HWW:1999} showed that a good fit could be obtained to the multi-color light curves of SN~1998bw near maximum with an asymmetric model.  This model required somewhat higher kinetic energy and \\ni\\ mass, $2\\times10^{51}$ erg and 0.2 \\Msun, respectively, but not sufficiently large values to warrant defining a new category of ``hypernovae.\" Furthermore, models of jet-induced supernovae \\citet{Khokhlov:1999} showed that asymmetric models would undergo circulation and inflow that would yield much smaller minimum velocities than the corresponding spherical models.  These small minimum velocities were revealed by late time observations of SN~1998bw (Patat et al. 2001). The slow moving matter will, in principle, allow for greater trapping of \\grs\\ at later times and hence flatter late-time light curves requiring less \\ni, all else being equal. \\citet{Maeda:2002} also conclude that asymmetric explosions can lead to slower moving matter near the center, but they only explored very high mass carbon/oxygen cores in the context of SN~1998bw, and hence concluded once again that high kinetic energy is required to match SN~1998bw.  Estimates of the amount of \\ni\\ required to power the light curve of SN~1998bw have also evolved with time.  Based on the first 123 days, \\citet{McKenzie:1999} estimated a minimum of 0.2 \\Msun.  While this number agreed with the estimate of \\citet{HWW:1999}, it was considerably less than the estimates based on the first spherically-symmetric light curve models.  A later, very careful analysis of the late-time light curve, especially HST photometry at 800 to 1000 d, by \\citet{Sollerman:2002} showed that a simple model could fit the light curve with only 0.3 \\Msun\\ of \\ni.  More sophisticated models allowing for the slowly expanding and efficiently trapping inner layers expected from asymmetric models might very well do the job with less \\ni.  There is, in any case, little rationale for the early large estimates of \\ni\\ mass based on spherical models.  In this context, we note that the ejected \\ni\\ mass for normal SN~II is reported to range as high as 0.3 \\Msun\\ \\citep{Schmidt:1994}. \\citet{Germany:2000} report an implied \\ni\\ mass of 2.6 \\Msun\\ for SN~1997cy, but this number is so discrepant with other determinations of the \\ni\\ mass that we are tempted to assume a substantial proportion of the luminosity of this event arise from circumstellar interaction. More events like SN~1997cy will have to be discovered to clarify this situation.  If one removes the seeming excessively large nickel mass and can, in principle, account for the observed optical luminosity of SN~1998bw by invoking the angle-dependent luminosity expected from asymmetric models, and can account for the large photospheric velocities by a judicious choice of ejecta mass and viewing angle, then one still has the active possibility that SN~1998bw was an event at the high end of the ``normal\" range of SN~Ic supernovae with twice the kinetic energy and twice the nickel mass of ``normal\" SN~Ic.  One does not necessarily need to define a whole new class of ``hypernovae\" to account for it.  Similar arguments can account for events that are perceived to be intermediate between SN~1998bw and SN~1994I, such as SN1998ef.  Even if one does not accept that one can account for the observations of SN~1998bw, SN~1998ef, and SN~2002ap in terms of mild variations and rather extreme asymmetries, the polarimetry shows that SN~1998bw and SN~2002ap, at least, were asymmetric and this must be taken into account. The \\grb\\ community recognized the key role that collimation plays in \\grbs\\ by adopting the language of ``isotropic equivalent energy\" in appropriate contexts when spherically symmetric models were applied to asymmetric situations.  Perhaps it is time for the supernova community to do that as well and learn to speak of isotropic equivalent energies, ejecta masses, and nickel masses when spherically-symmetric models are inappropriately applied to asymmetric situations."
    },
    "0206/astro-ph0206453_arXiv.txt": {
        "abstract": "We present a status report on our high speed photometric survey of faint Cataclysmic Variables, which is concentrating on old novae. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206447_arXiv.txt": {
        "abstract": "The \\HI\\ Parkes\\footnote{The Parkes telescope is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.} All-Sky Survey (HIPASS) is a blind 21-cm survey for extragalactic neutral hydrogen, covering the whole southern sky. The HIPASS Bright Galaxy Catalog (BGC; Koribalski et al. 2002) is a subset of HIPASS and contains the 1000 \\HI-brightest (peak flux density) galaxies. Here we present the 138 HIPASS BGC galaxies, which had no redshift measured prior to the Parkes multibeam \\HI\\ surveys.  Of the 138 galaxies, 87 are newly cataloged. Newly cataloged is defined as no optical (or infrared) counterpart in the NASA/IPAC Extragalactic Database. Using the Digitized Sky Survey we identify optical counterparts for almost half of the newly cataloged galaxies, which are typically of irregular or magellanic morphological type. Several \\HI\\ sources appear to be associated with compact groups or pairs of galaxies rather than an individual galaxy. The majority (57) of the newly cataloged galaxies lie within ten degrees of the Galactic Plane and are missing from optical surveys due to confusion with stars or dust extinction. This sample also includes newly cataloged galaxies first discovered in the \\HI\\ shallow survey of the Zone-of-Avoidance (Henning et al. 2000). The other 30 newly cataloged galaxies escaped detection due to their low surface brightness or optical compactness. Only one of these, HIPASS J0546--68, has no obvious optical counterpart as it is obscured by the Large Magellanic Cloud. We find that the newly cataloged galaxies with \\bgt\\ are generally lower in \\HI\\ mass and narrower in velocity width compared with the total HIPASS BGC. In contrast, newly cataloged galaxies behind the Milky Way are found to be statistically similar to the entire HIPASS BGC. In addition to these galaxies, the HIPASS BGC contains four previously unknown \\HI\\ clouds. ",
        "introduction": "\\label{sect:intro} \\setcounter{footnote}{1} The 21-cm line of neutral hydrogen (\\HI) is unique as it can probe regions of the sky where no stars have (yet) formed (see Schneider 1996). Within individual galaxies, \\HI\\ is frequently found well outside the optical radius (e.g. Meurer et al. 1996; Salpeter \\& Hoffman 1996), and many tidal tails or bridges between galaxies are only detected in \\HI\\ (see e.g. Koribalski 1996; Ryder et al. 2001). Until now, the majority of \\HI\\ observations were made of objects that had first been identified in the optical (or lately the infrared), thus imposing \\HI\\ selection effects on top of already existing optical selection effects. Important \\HI\\ structures like the Leo ring (Schneider 1989) and the Virgo cloud (Giovanelli \\& Haynes 1989) were discovered by accident and indicate the enormous potential for discovery in an untargeted \\HI\\ survey.  The sky has been extensively surveyed for galaxies at optical wavelengths (e.g. Lauberts 1982) but severe limitations remain, mainly due to the foreground extinction of the Milky Way (which affects $\\sim$25\\% of the sky, see e.g. Kraan-Korteweg \\& Lahav 2000).  In many optical catalogs, including the input catalogs for optical redshift surveys, low-surface brightness (LSB) galaxies are easily missed and galaxies with diameters less than $\\sim$1\\arcmin\\ are often misclassified as stars. For example, all objects with brightness less than 1.15 times the sky and objects classified as stars were excluded from the input catalog of the Las Campanas Redshift Survey (Shectman et al. 1996). To supplement the galaxy catalogs, targeted searches for LSB galaxies (Schneider et al. 1990, 1992; Impey et al. 1996, 2001; Morshidi-Esslinger et al. 1999; Cabanela \\& Dickey 2000) and dwarf galaxies (Karachentseva \\& Karachentsev 1998; Drinkwater et al. 1999; Drinkwater et al. 2000) as well as deep optical searches for galaxies behind the southern Milky Way (Woudt \\& Kraan-Korteweg 2001) are being carried out. In the infrared less than $10\\%$ of the sky is affected by foreground extinction, surveys like the Two Micron All Sky Survey (2MASS, Jarrett et al. 2000) and the Deep Near-Infrared Survey of the southern sky (DENIS, Epchtein et al.  1997) are now cataloging large numbers of galaxies.  In contrast \\HI\\ emission is {\\it not} affected by extinction and enables us to identify many previously hidden galaxies. In addition, \\HI\\ can easily be detected in LSB and late-type dwarf galaxies which are generally gas-rich (Impey \\& Bothun 1997).  \\HI\\ surveys complement optical galaxy catalogs and substantially improve the census of galaxies and measurement of the \\HI\\ content of the local Universe. \\HI\\ surveys also clarify voids by placing reliable upper limits on the mass of objects in these regions.  Furthermore, there are some components of galaxies that have so far only been discovered in \\HI, e.g. high-velocity clouds (Putman et al. 2002), tidal \\HI\\ clouds (e.g., HIPASS J0731--69, Ryder et al.  2001) and other nearby \\HI\\ clouds (see e.g. Kilborn et al. 2000). Elliptical galaxies, which are typically \\HI-poor, are the main component missing from \\HI\\ surveys (see e.g. Sanders, 1980; Knapp, Turner \\& Cunniffe, 1985)  The current view of the large-scale structure in the Local Universe, with its filaments and voids, is based almost entirely on optical observations of high-luminosity galaxies. This view is highly selective and it will be interesting to see how large-scale surveys for extragalactic neutral hydrogen affect the current picture. Until recently, the Arecibo Dual-Beam Survey (ADBS) and the deeper Arecibo \\HI\\ Strip Survey (AHISS) were the largest blind 21-cm surveys, covering areas of 430 and 65 deg$^2$ respectively. Rosenberg \\& Schneider (ADBS, 2000) detected 265 galaxies, of which 81 were uncataloged, whereas Zwaan et al. (AHISS, 1997; see also Zwaan 2000) detected 66 galaxies, half of which were uncataloged. With the advent of a 21-cm multibeam system at the 64-m Parkes telescope (see Staveley-Smith et al. 1996) as well as new observing and data reduction software (Barnes et al.  2001), much larger and deeper surveys are now possible. The \\HI\\ Parkes All-Sky Survey (HIPASS, see e.g. Koribalski 2002) is the largest 21-cm survey for neutral hydrogen to date, covering the whole southern sky. With these surveys, extragalactic \\HI\\ astronomy no longer depends entirely on observations at other wavelengths.  There is a large potential for detecting invisible \\HI\\ clouds and uncataloged galaxies with unusual properties in HIPASS. We expect to find many uncataloged galaxies, either hidden behind the Milky Way with \\HI\\ properties similar to the overall galaxy population or missed due to optical/infrared selection criteria. The former are important for the completion of our picture of the local large-scale structure as they bridge previously known structures that are optically intercepted by the Galactic Plane (e.g. Henning et al. 2000; Sharpe et al.  2001). The latter are equally important to enhance the completeness of galaxy catalogs across all morphological types.  Subsets of HIPASS within particular regions of the sky have already been analysed. In each of these regions uncataloged galaxies were discovered, many of which are also in our sample. Five uncataloged galaxies have been found in the Centaurus\\,A group (Banks et al. 1999). The South Celestial Cap (SCC) region of the sky ($\\delta\\ < -62\\degr$) has been studied extensively by Kilborn et al. (2002; see also Kilborn 2001), who found 114 uncataloged galaxies (out of 536 galaxies in total). Banks et al. (1999) and Kilborn et al. (2002) searched the HIPASS data to full sensitivity, so only some of their galaxies will appear in the HIPASS Bright Galaxy Catalog (Koribalski et al. 2002). On-going analysis of the full-sensitivity HIPASS data over the total survey area is expected to reveal many more uncataloged galaxies.  Henning et al. (2000) searched the \\HI\\ Zone-of-Avoidance shallow survey (HIZSS; $212\\degr\\ \\leq\\ l \\leq 36\\degr, |b| \\leq\\ 5\\degr$) and found 110 \\HI\\ sources, 67 of which had no published optical counterpart (see also Staveley-Smith et al. 1998). An \\HI\\ survey of the Great Attractor region ($l$ = 300\\degr\\ to 332\\degr, $|b| < 5$\\degr) by Staveley-Smith et al. (2000) has so far revealed 305 galaxies, most of which were previously unknown (see also Juraszek et al. 2000). Complete analysis of deep Zone-of-Avoidance \\HI\\ data is under way (Staveley-Smith et al., in prep.).  The \\HI\\ properties of all 1000 galaxies in the HIPASS Bright Galaxy Catalog are presented by Koribalski et al. (2002). The optical properties of all previously catalogued galaxies in the HIPASS BGC are analysed by Jerjen et al. (2002).  And the \\HI\\ mass function for the HIPASS BGC will be discussed by Zwaan et al. (2002). Here we present the \\HI\\ properties of 138 galaxies from the HIPASS BGC without velocity measurements prior to the Parkes multibeam surveys; 87 of these galaxies are newly cataloged -- that is they do not have a cataloged optical (or infrared) counterpart listed in the NASA/IPAC Extragalactic Database (NED). The number distribution of the BGC galaxies presented in this paper are summarised in Table~\\ref{tab:numbers}. In Section~2 we briefly describe the observations and the HIPASS BGC selection criteria as well as the method for optical (and infrared) identifications. In Section~3 we compare the \\HI\\ properties of newly cataloged galaxies with low and high absolute Galactic latitudes. Section~4 contains the conclusions. A short description of all the newly cataloged galaxies with identified optical counterparts is given in the Appendix. ",
        "conclusions": "\\label{sect:con}  A blind \\HI\\ survey such as HIPASS provides a view of the local Universe free from optical selection effects. Although the potential for detecting previously unknown \\HI\\ structures is high, we do not find any invisible \\HI\\ clouds not gravitationally bound to any stellar system in the HIPASS Bright Galaxy Catalog (BGC). The four indentified \\HI\\ clouds are most likely associated with Magellanic debris or other visible galaxies (The NGC 2442 group in the case of J0731--69).  This can place important upper limits on the contribution of \\HI\\ gas, not associated with galaxies, to the local baryon density. The HIPASS BGC has improved the census of galaxies in the local Universe by detecting galaxies behind the Milky Way. Furthermore, over the whole sky, we have easily detected galaxies missed in traditional optical surveys due to low surface brightness or misclassification as stars.  There are 87 newly cataloged galaxies in the HIPASS BGC, an additional 51 galaxies had no redshift measurement prior to the Parkes \\HI\\ multibeam surveys. The majority (57) of the newly cataloged galaxies lie behind the Milky Way (\\blt) and are missing from optical catalogs due to confusion or dust extinction. Optical counterparts are found in the Digitized Sky Survey for only 14 of these galaxies. Statistically, these 57 galaxies are found to have a similar \\HI\\ mass distribution and velocity widths to the entire HIPASS BGC.  All the newly cataloged galaxies with high absolute Galactic latitudes (30) have a candidate optical counterpart(s) with morphologies ranging from late-type spiral to irregular, including four with multiple optical counterparts. The exception is HIPASS J0546--68, which lies behind the LMC and has no visible optical counterpart. The characteristic surface brightness of these galaxies is extreme, either diffuse low surface brightness or compact high surface brightness. Although these galaxies are \\HI-rich, they are not high in \\HI\\ mass. The newly cataloged galaxies with \\bgt\\ on average have a lower \\HI\\ mass (median log(\\mhi/\\msun) = 8.7) and narrower velocity width (mean \\wfi\\ = 64 \\kms) than \\HI\\ selected galaxies with optically cataloged counterparts."
    },
    "0206/astro-ph0206392_arXiv.txt": {
        "abstract": "We show that the collapsar model of gamma-ray bursts results in a series of successive shocks and rarefaction waves propagating in the ``cork\" of stellar material being pushed ahead of the jet, as it emerges from the massive stellar progenitor. Our results are derived from analytical calculations assuming a hot, ultrarelativistic 1-D flow with an initial Lorentz factor $\\Gamma_j\\sim 100$. The shocks result in a series of characteristic, increasingly shorter and harder thermal X-ray pulses, as well as a non-thermal $\\gamma$-ray pulse, which precede the usual non-thermal MeV $\\gamma$-rays. We consider jets escaping from both compact (CO or He dwarf) and blue supergiant stellar progenitors. The period, fluence and hardness of the pulses serves as a diagnostic for the size and density of the outer envelope of the progenitor star. ",
        "introduction": "The non-thermal radiation from shocks in a relativistic fireball jet after emerging from a collapsing massive stellar system (collapsar) is the leading theoretical explanation for the family of ``long\" gamma-ray bursts (GRB) with $\\gamma$-ray durations longer than about $t_b\\sim 10$ s (Woosley 1993, \\Pacz 1998; see \\Mesz 2002 for a review). The jet is thought to arise from a brief accretion episode of a rotating debris torus around a central black hole formed as the stellar core collapses. The nature and history of the massive star, however, is so far largely a matter of hypothesis and calculation, and alternative stellar progenitor systems are conceivable, both for the long and especially for the shorter bursts.  The duration of the TeV neutrino burst expected from shocks in the jet before it emerges is a possible diagnostic for the pre-emergence jet history and for the dimensions or column density of the progenitor stellar system (\\Mesz \\& Waxman 2001). Previous analytical work on shock and jet emergence has been done by Matzner \\& McKee 1999, \\Mesz \\& Rees 2001, Matzner 2002, and numerical calculations of relativistic jets have been done by Marti, M\\\"uller, Font, Iba\\~nez \\& Marquina, 1997; Aloy, M\\\"uller, Iba\\~nez, Marti \\& MacFadyen, 2000; and Zhang, Woosley \\& McFadyen, 2002. Here we concentrate on possible photon signatures of the jet as it emerges, which precedes the usual gamma-ray emission and the subsequent longer wavelength afterglow. This has been considered by Ramirez-Ruiz, McFadyen \\& Lazzati (2001), who infer an X-ray precursor to the burst from the portion of the stellar envelope shocked by the jet. Here we investigate in some detail, using analytical 1-D methods in the ultrarelativistic limit, the jet emergence process and the shock heating and expansion of the plug of stellar material propelled by it, resembling an ejected cork.  In the initial stage of ejection the cork is optically very thick to scattering, and it experiences successive shocks and rarefactions as the optical depth decreases during its overall expansion. This leads to an X-ray photon precursor to the usual GRB emission, consisting of a pattern with an initial rise and decay followed by a brighter and harder peak, possibly repeated a few times with decreasing period. The characteristic timescale of the initial decay and the ratio of its amplitude to that of the peak differ substantially depending on the density and extent of the envelope, thus providing a potential diagnostic for the pre-burst stellar configuration. These precursor patterns should be detectable with instruments to be flown in the next few years. ",
        "conclusions": "\\label{sec:disc}  We have shown that the emergence of a jet in a collapsar model of GRB leads to a series of successive, increasingly shorter and harder thermal X-ray pulses. These are caused by successive cycles of shock waves and rarefaction waves reaching the outer and inner ends of the ``cork\" of stellar material being pushed ahead of the jet as it passes beyond the boundary of the stellar envelope. The Lorentz factor of the main  bulk of the cork increases with each succeeding and increasingly relativistic shock that goes through it. A small fraction of the cork mass may be accelerated to larger Lorentz factors and run ahead of the cork, without changing significantly the development of the thermal pulses from the main bulk of the cork.  The expansion and acceleration of the main bulk of the cork lead to a decrease of the cork scattering optical depth, photon diffusion times and expansion times in the observer frame for the successive shocks. As a result the successive shocks have increasingly higher  bolometric fluences and shorter durations, as well as harder peak photon energies. However, for observations in a fixed low energy X-ray band, e.g 2 keV, the specific fluence in erg/cm$^2$/keV would be reduced by a Rayleigh-Jeans factor $(E_{obs}/E_{pk})^2$ and would appear almost constant or slowly declining as the pulse number increases. On the other hand, in a hard X-ray detector, e.g. $\\simg 20$ keV, the successive pulses would show a steeply increasing fluence.  The results discussed here are based on analytical 1-D calculations in the relativistic limit. As argued in \\S~2 and in \\S~3, the approximations inherent in such an analytical treatment are reasonable for the range of parameters considered, but the flows considered are quite complex, and definitive results must await confirmation from more accurate numerical simulations. Current 2-D relativistic collapsar jet simulations (e.g. Aloy \\etal 2000, Zhang \\etal, 2002) do not so far have the dynamic range and resolution necessary to distinguish multiple shocks and rarefaction waves upon emergence of the jet, but provide valuable informations about the overall jet dynamics. The approximate analytical treatment presented here fulfill in the meantime an exploratory role, and provide insights both for future numerical and observational developments.  The temporal structure discussed provides distinct features, which should help to identify the phenomenon. We have argued (\\S \\ref{sec:jet-prop}) that since jets are likely to accelerate to significantly relativistic velocities as they enter a much lower density H-envelope, the crossing time for a blue supergiant envelope may be only slightly longer than the core crossing time $t_{core}\\sim 30r_{11}$ s for He or CO dwarf stars. The initial rise-time of the first pulse leading to the peak flux, given by the angular smearing time, is $t_{pk}^{obs}\\sim [500,~20,~2]$ ms for a blue supergiant (H) star, a He star and a CO star, occurring while the cork has not yet had time to expand. This is followed by a decay $f\\propto t^{-1/2}$ lasting an expansion time $t_{exp}^{obs}\\sim [2,~0.15,~0.025]$ s for the [H, He, Co] cases. Thereafter the flux drops as $f\\propto t^{-2}$. The subsequent pulses have similar rise and decay times, given approximately by (increasingly shorter) observer-frame expansion times.  For stars of greater compactness (smaller radius) the characteristic bolometric fluxes, timescales and fluences of each pulse are lower and shorter while the peak temperatures are higher. The typical fluences from a Hubble distance $d\\sim 10^{28}$ cm range from ${\\cal F} \\sim [5\\times 10^{-9},7\\times 10^{-12},5\\times 10^{-14}]$  erg/cm$^2$ for the first pulse, to ${\\cal F} \\sim [8\\times 10^{-7},5\\times 10^{-10},2\\times 10^{-12}]$ erg/cm$^2$ for the third pulse, while the peak temperatures range from $T_{pk}^{obs}\\sim [1.6,~12,~50]$ keV for the first pulse, to $T_{pk}^{obs}\\sim [14,~52,~110]$ keV for the third pulse in [H, He, CO] stars respectively.  After the third shock, the cork will have reached Lorentz factors comparable to the observationally inferred Lorentz factors in bursts ($\\Gamma \\sim 100)$. This shock may become collisionless, leading also to an additional non-thermal component of energy and duration comparable to the thermal third pulse, with characteristic photon energy $E_\\gamma \\sim [8,~45,~140]$ MeV for [H, He, CO] stars.  With the fluences calculated here it is apparent that it would be easier to detect the X-ray pulses from extended GRB progenitors, such as blue supergiant (H) stars.  The more compact progenitors (He or CO stars) would have lower fluences and be harder to detect, except for nearby objects, in which case they would be distinguishable from blue supergiant progenitors through their shorter peak durations and harder spectra. The characteristic precursor thermal pulses, including the third thermal and non-thermal pulses described here, would precede the conventional MeV range emission by an appreciable timescale $\\sim t_{exp}^{obs} \\sim 2.2 \\L52^{-1/2}\\theta_{-1}^{3/2} r_{12.5}^{3/2}\\rho_{-7}^{1/4}$ s for blue supergiant (H) progenitors, or by shorter timescales $\\sim 0.15,~0.025$ s for He or CO progenitors.  It is conventionally assumed that the conditions for the usual successful GRB to occur is that the central engine powers the jet for timescales $t_j$ longer than the stellar crossing time $t_{core} \\sim 30r_{11}$ s. However, there may be cases where the jet just breaks through the outer envelope of the star (whether it be a He, CO or H BSG star), and does not live much longer afterwards (e.g. $t_j \\siml 30$ s). In this case, electromagnetic detectors will see only the thermal X-ray pulses and the hard non-thermal gamma-ray pulse ($\\simg 10$ MeV, \\S \\ref{sec:non-thermal}), which last seconds or less. These would be a new class of objects, perhaps related to X-ray flash bursts and/or to short GRBs, especially if such events are more common, and their distance is smaller.  The characteristic timescales as well as the ratios of the amplitudes and hardnesses of the pulses differ substantially, depending on the density and extent of the stellar envelope, which provides a potentially valuable diagnostic for pinning down the conditions in the pre-burst stellar progenitor.  Instruments to be flown in the next several years, such as Swift, GLAST, AGILE and others, should be able to detect such precursor signals in some bursts, and provide valuable clues concerning the GRB mechanism and their progenitors."
    },
    "0206/astro-ph0206501_arXiv.txt": {
        "abstract": "We present a new, model-independent approach for measuring non-Gaussianity of the Cosmic Microwave Background (CMB) anisotropy pattern. Our approach is based on the empirical distribution function of the normalized spherical harmonic expansion coefficients $a_{\\ell m}$ of a nearly full-sky CMB map, like the ones expected from forthcoming satellite experiments. Using a set of Kolmogorov-Smirnov type tests, we check for Gaussianity and independency of the $a_{\\ell m}$. We test the method on two non-Gaussian toy-models of the CMB, one generated in spherical harmonic space and one in pixel (real) space. We also provide some rigorous results, possibly of independent interest, on the exact distribution of the spherical harmonic coefficients normalized by an estimated angular power spectrum. ",
        "introduction": "\\label{sect:intro}  Temperature fluctuations in the CMB are an invaluable tool to constraint cosmological models and the process of structure formation in the universe. According to the standard version of the most popular theory of the very early universe, the so-called theory of inflation, the density fluctuations at primordial epochs should be Gaussian or very close to Gaussian distributed (see the reviews \\cite{narlikar,watson}). This implies that the temperature fluctuations in the CMB as observed today should also be close to Gaussian, because they are related by linear theory to fluctuations in the early universe. However, the details of the inflationary scenario are still rather unclear. Newer, more sophisticated, versions of inflation predict small deviations from Gaussianity \\cite{nongi1,nongi2,nongi3,nongi4,nongi5,nongi6}. Detecting this non-Gaussianity in the CMB would thus be important for understanding the physics of the very early universe. There are also other possible mechanisms for the creation of  non-Gaussianity in the CMB. If the universe has undergone a phase transition at early times, this could have given rise to topological defects (See Ref.\\cite{strings} for a review). These defects would show up as non-Gaussian features in the CMB temperature fluctuation field.\\\\  With the high resolution data from the satellite missions MAP\\footnote{http://map.gsfc.nasa.gov/} and Planck Surveyor\\footnote{http://astro.estec.esa.nl/SA-general/Projects/Planck/} one could be able to detect possible deviations from non-Gaussianity in the CMB. If this is indeed detected it would have a big impact on our understanding of the physics of the early universe. For these reasons, the search for procedures to test for Gaussianity in high resolution data has recently drawn an enormous amount of attention in the CMB literature. A number of methods were proposed, many based upon topological properties of spherical Gaussian fields: the behavior of Minkowski functionals \\cite{novikov,gott}, temperature correlation functions \\cite{eriksen}, the peak to peak correlation function \\cite{heavens}, skewness and kurtosis of the temperature field \\cite{vittorio} and local curvature properties of Gaussian and chi-squared fields \\cite{dore}. Other works have focussed on harmonic space approaches: analysis of the bispectrum and its normalized version \\cite{phillips,komatsu} and the bispectrum in the flat sky approximation \\cite{win}. The explicit form of the trispectrum for CMB data was derived in \\cite{hu,kunz}. Applications to COBE, Maxima and Boomerang data have also drawn enormous attention and raised wide debate \\cite{boom,polenta,cobeng1,cobeng2}.\\\\  Another reason to look for non-Gaussianities in CMB data, is to detect effects of systematic errors in the CMB map. One of these systematic effects can be stripes from the $1/f$ noise in the detectors which have not been properly removed. Another systematics which could induce non-physical non-Gaussianities in the map could be the effect of straylight contamination from the galaxy or distortions of the main beam caused by the optics of the telescope. Finally astrophysical foregrounds like synchrotron emission, thermal dust emission or free-free emission from the galaxy and other extra galactic sources could be wrongly interpreted as a physical non-Gaussian signal. It is important that one uses data from simulated experiments to check whether these systematic effects induce non-physical non-Gaussian signatures in the CMB map.\\\\     Our purpose here is to propose a new procedure to detect non-Gaussianity in harmonic space. More precisely, let $T(\\theta ,\\varphi )$ denote the CMB fluctuations field, which we assume, as always, to be homogeneous and isotropic, for $0<\\theta \\leq \\pi ,$ $0<\\varphi \\leq 2\\pi $. Assuming that $% T(\\theta ,\\varphi )$ has zero mean and finite second moments, it is well-known that the following spectral representation holds: \\begin{equation*} T(\\theta ,\\varphi )=\\sum_{\\ell=1}^{\\infty }\\sum_{m=-\\ell}^{\\ell}a_{\\ell m}Y_{\\ell m}(\\theta ,\\varphi )\\text{ ,} \\end{equation*}% where $Y_{\\ell m}(\\theta ,\\varphi )$ denotes the spherical harmonics. The random coefficients (amplitudes) $\\left\\{ a_{\\ell m}\\right\\} $ have zero-mean with variance $\\langle|a_{\\ell m}|^{2}\\rangle=C_{\\ell}$. They are uncorrelated over $\\ell$ and $|m|$: $\\langle a_{\\ell m}a_{\\ell^{\\prime }m^{\\prime }}^{\\ast }\\rangle=C_{\\ell}\\delta _{\\ell}^{\\ell^{\\prime }}\\delta _{m}^{m^{\\prime }}$, and $a_{\\ell,m}=a_{\\ell,-m}^{% \\ast }$. The sequence $\\left\\{ C_{\\ell}\\right\\} $ denotes the angular power spectrum of the random field and the asterisk complex conjugation. Furthermore, if $T(\\theta ,\\varphi )$ is Gaussian, the $\\left\\{ a_{\\ell m}\\right\\} $ have a complex Gaussian distribution. Upon observing $% T(\\theta ,\\varphi )$ on the full sky, the random coefficients can be obtained through the inversion formula% \\begin{equation} a_{\\ell m}=\\int_{-\\pi }^{\\pi }\\int_{0}^{\\pi }T(\\theta ,\\varphi )Y_{\\ell m}^{\\ast }(\\theta ,\\varphi )\\sin \\theta d\\theta d\\varphi \\text{ , }m=0,\\pm 1,...,\\pm l% \\text{ },\\text{ }l=1,2,...\\text{ .}  \\label{alm} \\end{equation}% Our purpose is to study the empirical distribution function for the $\\left\\{ a_{\\ell m}\\right\\} ,$ and to use these results to implement tests for non-Gaussianity in harmonic space. We shall assume that the angular power spectrum is unknown, and the sequence $\\left\\{ C_{\\ell}\\right\\} $ estimated from the data; as we show below, this has a nonnegligible effect on the behavior of the test, no matter how good is the resolution of the experiment% $.$ The plan of the paper is as follows: the procedure we advocate is described in Sections \\ref{sect:univar} and \\ref{sect:multivar}; Section \\ref{sect:empres} presents some empirical results, whereas Section \\ref{sect:disc} is devoted to discussion and to directions for further work. Some mathematical results are collected in the Appendix. ",
        "conclusions": "\\label{sect:disc}  We believe the approach advocated here may enjoy some advantages over existing methods, and we view it as complementary to geometric approaches in pixel space (for instance, methods based an Minkowski functionals, local curvature, or other topological properties). Our proposal allows for a rigorous asymptotic theory; it is completely model free; it provides information not only on the existence of non-Gaussianity, but also on its location in the space of multipoles; the effect of estimated parameters is carefully accounted for; given that the asymptotic behavior of the field $% K_{L}$ has been thoroughly investigated, many other procedures, further than those we considered here (Kolmogorov-Smirnov and Cramer-Von Mises tests) can be immediately implemented; the extension to an arbitrary number of rows is in principle straightforward (although computationally burdensome), whereas for instance the explicit form of higher order cumulant spectra needs to be derived analytically in a case by case fashion. Furthermore, our analysis of the distributional properties of the spherical harmonic coefficients may have some independent interest in other areas of CMB investigation.  Our approach is clearly related to the analysis of higher-order cumulant spectra (such as the bispectrum and trispectrum) in harmonic space. Although the bi- and trispectrum have been very widely used in empirical work, their power properties against a variety of non-Gaussian models do not seem to have been very much investigated. We conjecture that our procedure may enjoy better power properties than higher order spectra in a number of circumstances; heuristically, the bispectrum and the trispectrum search for non-Gaussian features on the $a_{\\ell m}$ by focusing essentially on their skewness and kurtosis, whereas the method we advocate here probe their whole multivariate distribution. The mutual interplay between these different methodologies may lead to improvements in both directions: for instance, it seems possible to model the bispectrum and trispectrum evaluated at different multipoles as processes indexed by some $r$, much the same way as we did here to combine the information from empirical processes $% G_{\\ell_{1}...\\ell_{k}}$ into a single statistic $K_{L}.$ This may allow for a more rigorous analysis in the aggregate, in order to understand whether a single or a few high values are to be considered significant, over a set of $% L$ statistics. On the other hand, incorporation of some selection rules (e.g., the Wigner's $3j$ coefficients) into our analysis may help us to exploit the isotropic nature of the field to probe non-Gaussianity more efficiently.  In this paper we have tested the method for two non-Gaussian models. One which was created in spherical harmonic space and the other which was created in pixel space. For the first model, we detect non--Gaussianity at a $2\\sigma$ level (about $50\\%$ of the times) with the univariate test even when the test model contains only $5.9\\%$ non-Gaussianity. In this case the map is very similar to a Gaussian map. Our second non-Gaussian test model was generated in pixel space. In this case we had $2\\sigma$ detections (about $50\\%$ of the times) with $15.5\\%$ non-Gaussianity. The map shows that this kind of non-Gaussianity may be more easily detected in pixel space. When taking into account realistic effects like detector noise and galactic cuts, the method has so far shown promising results. This will be discussed further together with tests on realistic non-Gaussian maps in \\cite{paper2}."
    },
    "0206/astro-ph0206051_arXiv.txt": {
        "abstract": "{ This paper presents the analysis of Doppler p-mode observations of the G2V star $\\alpha$~Cen~A obtained with the spectrograph CORALIE in May 2001. Thirteen nights of observations have made it possible to collect 1850 radial velocity measurements with a standard deviation of about 1.5 m\\,s$^{-1}$. Twenty-eight oscillation modes have been identified in the power spectrum between 1.8 and 2.9~mHz with amplitudes in the range 12 to 44 cm\\,s$^{-1}$. The average large and small spacing are respectively equal to 105.5 and 5.6~$\\mu$Hz. A comparison with stellar models of $\\alpha$~Cen~A is presented. ",
        "introduction": "The lack of observational constraints leads to serious uncertainties in the modeling of stellar interiors and stellar evolution. The measurement and characterization of oscillation modes in solar-like stars is an ideal tool to test models of stellar inner structure and theories of stellar evolution.  The discovery of propagating sound waves, also called p-modes, in the Sun by Leighton et al. (\\cite{leighton62}) and their interpretation by Ulrich (\\cite{ulrich70}), has opened a new area in stellar physics. Frequency and amplitude of each oscillation mode depend on the physical conditions prevailing in the layers crossed by the waves and provide a powerful seismological tool. Helioseismology led to major revisions in the ``standard model'' of the Sun and provided for instance measures of the inner rotation of the Sun, the size of the convective zone and the structure of the external layers.  Solar-like oscillation modes generate periodic motions of the stellar surface with periods in the range 3 - 30 minutes but with extremely small amplitudes. Essentially two methods exist to detect such a motion: photometry and Doppler spectroscopy. In photometry, the oscillation amplitudes of solar-like stars are in the range 2 - 30~ppm while they are in the range 10 - 150~{\\cms} in radial velocity measurements. Photometric measurements made from the ground are strongly limited by scintillation noise. To reach the needed accuracy requires observations made from space. In contrast, Doppler ground-based measurements have recently shown their capability to detect oscillation modes in solar-like stars.  The first good evidence of excess power due to mode oscillations was obtained by Martic et al. (\\cite{martic99}) on the F5 subgiant Procyon. Bedding et al. (\\cite{bedding01}) obtained a quite similar excess power for the G2 subgiant $\\beta$ Hyi. These two results have independently been confirmed by our group based on observations made with the CORALIE spectrograph (Carrier et al. \\cite{carrier01a}, \\cite{carrier01b}).  A primary target for the search for p-mode oscillations is the solar twin $\\alpha$~Cen~A (HR5459). Several groups have already made thorough attempts to detect the signature of p-mode oscillations on this star. Two groups claimed mode detections with amplitudes 3.2\\,-\\,6.4 greater than solar (Gelly et al. \\cite{gelly86}; Pottasch et al. \\cite{pottasch92}). However, this was refuted by three other groups who obtained upper limits of mode amplitudes of 1.4\\,-\\,3 times solar (Brown \\& Gilliland \\cite{brown90}; Edmonds \\& Cram \\cite{edmonds95}; Kjeldsen et al. \\cite{kjeldsen99}). More recently Schou \\& Buzasi (\\cite{schou01}) made photometric observations of $\\alpha$~Cen~A with the WIRE spacecraft and reported a possible detection. The first unambiguous observation of p-modes in this star was recently made with the spectrograph CORALIE mounted on the 1.2-m Swiss telescope at La Silla Observatory (Bouchy \\& Carrier \\cite{BC01} (hereafter referred as BC01); Carrier et al. \\cite{carrier01c}). We present here a revised and extended analysis of the acoustic spectrum of this star and compare it with theoretical models. ",
        "conclusions": "Our observations of $\\alpha$~Cen~A yield a clear detection of p-mode oscillations. Several identifiable modes appear in the power spectrum between 1.8 and 2.9 mHz with an average large spacing of 105.5 $\\mu$Hz, an average small spacing of 5.6 $\\mu$Hz and an envelope amplitude of about 31 \\cms. These characteristics, in full agreement with the expected values, make it possible to put constraints on the physical parameters of this star.  Additional data with a higher signal-to-noise and a higher frequency resolution will make it possible to determine properly the suspected rotational splitting and the damping time of $\\alpha$~Cen~A p-modes.  This result, obtained with a small telescope, demonstrates the power of Doppler ground-based asteroseismology. Future spectrographs like HARPS (Queloz et al. \\cite{queloz01}; Bouchy \\cite{bouchy01}) are expected to conduct asteroseismological study on a large sample of solar-like stars, and to enlarge significantly our understanding of stellar physics."
    },
    "0206/astro-ph0206098_arXiv.txt": {
        "abstract": "We use a large set of state-of-the-art cosmological N-body simulations [$512^3$ particles] to study the intrinsic ellipticity correlation functions of halos. With the simulations of different resolutions, we find that the ellipticity correlations converge once the halos have more than 160 members. For halos with fewer members, the correlations are underestimated, and the underestimation amounts to a factor of 2 when the halos have only $20$ particles. After correcting for the resolution effects, we show that the ellipticity correlations of halos in the bigger box ($L=300\\mpc$) agree very well with those obtained in the smaller box ($L=100\\mpc$). Combining these results from the different simulation boxes, we present accurate fitting formulae for the ellipticity correlation function $c_{11}(r)$ and for the projected correlation functions $\\Sigma_{11}(r_p)$ and $\\Sigma_{22}(r_p)$ over three orders of magnitude in halo mass. The latter two functions are useful for predicting the contribution of the intrinsic correlations to deep lensing surveys. With reasonable assumptions for the redshift distribution of galaxies and for the mass of galaxies, we find that the intrinsic ellipticity correlation can contribute significantly not only to shallow surveys but also to deep surveys. Our results indicate that previous similar studies significantly underestimated this contribution for their limited simulation resolutions. ",
        "introduction": "Inhomogeneities of matter distribution in the Universe distort the images of distant galaxies gravitationally, a phenomenon called gravitational lensing. The lensing effect induces an ellipticity-ellipticity correlation of the galaxies on large scales, which is observable and can be used as a powerful tool to probe the large-scale dark matter distribution in the Universe (Miralda-Escude 1991; Blandford et al. 1991; Kaiser 1992; see Bartelmann \\& Schneider 2001 and Mellier 1999 for reviews).  Several groups have already detected the ellipticity correlation on scales from $0.4$ to $30$ arc-minutes for faint galaxies (Bacon et al. 2000, 2002; Hoekstra et al.  2002a,b; Kaiser et al. 2000; Maoli et al. 2001,; Rhodes et al.~2000, 2001; van Waerbeke et al. 2000, 2001; Wittman et al. 2000). If the source galaxies are randomly oriented without the lensing effect, that is, the intrinsic ellipticity correlation of the galaxies is negligible, the observed ellipticity correlation implies that the parameter $\\beta\\equiv \\Omega_0^{0.6}\\sigma_8$ is about 0.6, where $\\Omega_0$ is the density parameter, $\\sigma_8$ is the current rms linear density fluctuation in a top-hat sphere of $8\\mpc$, and $h$ is the Hubble constant in units of $100 {\\rm km~s^{-1} Mpc^{-1}}$.  There are however evidences both from theory and from observations that the shapes of galaxies are intrinsically correlated. In the theory, the large scale tidal field is expected to induce large-scale correlations in galaxy spins and in galaxy shapes (Lee \\& Pen 2000, 2001; Croft \\& Metzler 2000, hereafter CM00; Heavens et al. 2000, hereafter HRH00; Catelan et al. 2001; Catelan \\& Porciani 2001; Hui \\& Zhang 2002; Porciani et al. 2002). It is recently claimed that a large-scale alignment of galaxy spins has been detected in nearby galaxy catalogs with a high confidence (Pen et al. 2000, Lee \\& Pen 2002; Brown et al. 2002; Plionis 1994 for cluster shapes). While the intrinsic ellipticity correlation may be separated from the weak lensing signal in observations through measuring the $E$-mode and the $B$-mode correlations of the ellipticity (Crittenden et al. 2001, Mackey et al. 2002), applying this technique needs an accurate relation between the $E$-mode and $B$-mode correlations (Crittenden et al. 2001, 2002; Pen et al. 2002; Schneider et al. 2002; Hoekstra et al. 2002b). In this aspect, the current situations are far from satisfactory, since there are still considerable uncertainties both in the theoretical predictions and in the observational measurements of the intrinsic ellipticity correlations.  Using the numerical simulations of $256^3$ particles released by the Virgo Consortium, HRH00 and CM00 measured the ellipticity correlation for dark matter halos and discussed their results in the context of weak lensing measurements\\footnote{HRH00 also examined the correlation of the halo spins.}.  Assuming that the galaxies have the same intrinsic correlations as their host halos, they found that the intrinsic correlation of galaxy ellipticity could dominate over the lensing signal in shallow lensing surveys, and could contribute a non-negligible signal, as high as 20 percent, to deep lensing surveys like those that reported detections of the weak lensing effects. Their findings already have profound implications for interpreting the weak lensing experiments, but quantitatively speaking, there are still significant uncertainties in their results.  For examples, CM00 compared the ellipticity correlation $c_{11}(r)$ for halos of a mass $>1.4\\times 10^{12}\\msun$ between two simulations of boxsizes $L= 141\\mpc$ and $240\\mpc$, and found that $c_{11}(r)$ in the higher resolution simulation ($L=141\\mpc$) is $2\\sim 3$ times higher.  The ellipticity correlation $c_{11}(r)$ found in CM00 is also a factor of $2$ or more higher than that found in HRH00. Therefore, it is unclear how their selections of halos and simulation resolutions have affected their results.  In this {\\it Letter}, we will first quantify how the simulation resolution affects the determination of the ellipticity correlation using a set of state-of-the-art cosmological N-body simulations of $512^3$ particles (Jing \\& Suto 2002) and a set of lower-resolution simulations of $256^3$ particles (Jing 1998; Jing \\& Suto 1998). We will show that the minimum particle number $N_{\\rm min}$ for resolving the halo shapes in simulations is $160$, and the ellipticity correlation functions of halos with more than $N_{\\rm min}$ particles are well converged. We will also show that the ellipticity correlation is underestimated by a factor $2$ when halos of about twenty particles are used. Therefore, the predictions of CM00 and HRH00 for weak lensing surveys are significantly underestimated, since the halos of ten (HRH00) or twenty (CM00) particles were included in their studies (in order to have a sufficient number of halos for their analyses). Our results imply that the intrinsic ellipticity correlation may have a significant contribution even to deep weak lensing surveys. Our simulations are also large enough for accurately measuring the scale- and mass-dependences of the ellipticity correlation functions. Based on the simulation data, accurate fitting formulae are presented for these functions. These formulae can be used to predict the intrinsic ellipticity signals, including $E$-mode and $B$-mode contributions, in large-scale lensing surveys (Crittenden et al. 2001, 2002; Schneider et al. 2001; Pen et al. 2002).  In the next section, we will measure the ellipticity correlation functions in a set of N-body simulations, and will do the convergence test. In Section 3, we will present the projected ellipticity correlation functions, which are useful for predicting the intrinsic ellipticity contribution in weak lensing surveys once the radial (or redshift) distribution of source galaxies is known. Our results are summarized and discussed in Section 4. ",
        "conclusions": "We used a set of high-resolution cosmological N-body simulations to study the intrinsic ellipticity correlation functions of halos. With the simulations of different resolutions, we found that the ellipticity correlations converge once the halos have more than 160 members. For halos with fewer members, the correlations are underestimated and the underestimation amounts to a factor of 2 when the halos have only $20$ particles. After correcting for the resolution effects, we found that the ellipticity correlations of the halos in the bigger box ($L=300\\mpc$) agree very well with those obtained in the small box ($L=100\\mpc$). Combining these results from the different simulation boxes, we have presented accurate fitting formulae for the ellipticity correlation functions $c_{11}(r)$ and for the projected correlation functions $\\Sigma_{11}(r_p)$ and $\\Sigma_{22}(r_p)$ over a large range of halo mass (at least for $10^{10}\\le M_h \\le 10^{13}\\msun$). The latter two functions can be used to predict the contribution of the intrinsic correlations to deep lensing surveys. With reasonable assumptions for the redshift distribution of galaxies and for the mass of galaxies, we found that the intrinsic ellipticity correlation can contribute significantly not only to shallow surveys but also to deep surveys, if the galaxies have the same shapes as their host halos."
    },
    "0206/astro-ph0206267_arXiv.txt": {
        "abstract": "{ It is shown that the lens equation for a binary gravitational lens being a set of two coupled real fifth-order algebraic equations (equivalent to a single complex equation of the same order) can be reduced to a single real fifth-order algebraic equation, which provides a much simpler way to study lensing by binary objects. ",
        "introduction": "The gravitational lensing due to a binary system has attracted a lot of interest since the pioneering work by Schneider and Wei{\\ss} (1986). The lens equation adopted until now is a set of two coupled real fifth-order equations, equivalently a complex fifth-order equation (Witt 1990) which is based on a complex notation introduced by Bourassa, Kantowski and Norton (1973, 1975).  The number of images is classified by curves called caustics, on which the Jacobian of the lens mapping vanishes on a source plane. Caustics for two-point masses are investigated in detail and locations of caustics are clarified based on a set of two coupled real fifth-order equations as an application of catastrophe theory (Erdl and Schneider 1993) and on a complex formalism (Witt and Petters 1993), which is developed as an efficient method to compute microlensed light curves for point sources (Witt 1993): In the binary lensing, three images appear for a source outside the caustic, while five images are caused for a source inside the caustic. For a symmetric binary with two equal masses, the lens equation for a source on the symmetry axes of the binary becomes so simple that we can find the analytic solutions (Schneider and Wei{\\ss} 1986). In star-planet systems, the mass ratio of the binary is so small that we can find approximate solutions in general (Bozza 1999, Asada 2002). The approximate solutions are used to study the shift of the photocenter position by the astrometric microlensing (Asada 2002).  Nevertheless, it is quite difficult to solve these equations, since there are no well-established methods for solving coupled nonlinear equations numerically with sufficient accuracy (Press et al. 1988). We show that the lens equation can be reduced to a single master equation which is fifth-order in a real variable with real coefficients. As a consequence, the new formalism provides a much simpler way to study the binary lensing. ",
        "conclusions": "We have carefully reexamined the lens equation for a binary system in the polar coordinates. As a consequence, we have derived Eq. ($\\ref{theta2}$) for $\\tan\\phi$. After solving the equation, $r \\cos\\phi$ is determined by Eq. ($\\ref{rcos}$). Hence, the image position $(\\theta_x, \\theta_y)=(r \\cos\\phi, r \\cos\\phi \\tan\\phi) $ can be determined. Our formulation based on the one-dimensional equation ($\\ref{theta2}$) is significantly useful compared with previous two-dimensional treatments for which there are no well-established numerical methods (Press et al. 1988); the new formulation enables us to study the binary lensing more precisely with saving time and computer resources. For instance, it is effective in rapid and accurate light-curve fitting to microlensing events, in particular due to star-planet systems."
    },
    "0206/astro-ph0206034_arXiv.txt": {
        "abstract": "We have made the first direct interferometric proper motion measurements of the radio pulsar B1757$-$24, which sits at the tip of the ``beak'' of the putative ``Duck'' supernova remnant. The peculiar morphology of this radio complex has been used to argue alternately that the pulsar's space motion was either surprisingly high or surprisingly low.  In fact, we show that the pulsar's motion is so small that it and its associated nonthermal nebula G5.27$-$0.9 (the ``head'') are almost certainly unrelated to the much larger G5.4$-$1.2 (the ``wings''). ",
        "introduction": "The radio source colloquially known as the ``Duck'' (Fig.~1) consists of the large (35\\arcmin\\ diameter) fan-shaped nebula \\Gff\\ with edge brightening along the western edge (the ``wing''); the small (1\\arcmin\\ diameter) nebula \\Gft\\ a few arcminutes to the west (the ``head''); and a narrow (10\\arcsec\\ wide) extension reaching a further 30\\arcsec\\ west (the ``beak'').  At the very western tip is the 125~ms radio pulsar \\psr, with a characteristic spin-down age $\\tau=P/2\\dot P=15.5$~kyr.  Even before the localization of the pulsar or the measurement of its small apparent age, the Duck drew attention for its peculiar morphology \\citep{hb85,ckk+87}. The pulsar's youth and location near but outside the remnant only intensified the interest.  The large displacement of the pulsar from its putative birthplace has twice led to conclusions sufficiently remarkable to be widely reported in the press: first the apparent youth of the pulsar led to the conclusion it was traveling at $\\sim2000$~km/s \\citep{fk91,mkj+91}, then a surprisingly small limit on the motion of the pulsar bow shock led to the conclusion that the pulsar must be far older than its characteristic age \\citep[henceforth GF00]{gf00}.  The fundamental assumption on which these and scores of other publications depend is that the Duck is a coherent structure, not a chance line-of-sight juxtaposition of unrelated sources. This hypothesis allows several predictions: First, the distances to \\Gff, \\Gft, and \\psr\\ should all agree; second, the pulsar and remnant ages should agree; finally, the pulsar should be moving in a ballistic trajectory away from the supernova blast center.  Testing these predictions has been hard. The distance estimates from \\HI\\ absorption measurements and pulsar dispersion are not inconsistent, but are poorly constrained; the age estimates, initially in excellent agreement, have proven inconsistent with the proper motion limits; and the proper motion inferred from the bow shock is in the wrong direction.  Our goal was a direct measurement of the pulsar motion, to help elucidate the nature of the Duck. ",
        "conclusions": "We believe that the evidence linking \\psr\\ to the remnant \\Gff\\ is weak, consisting mainly of the relatively close projected proximity of the sources together with the unusual edge-brightening of the remnant. We have suggested that the former is not particularly surprising in this crowded part of the Galaxy, and the latter may be due to the brightening of the edge moving into the denser medium along the Galactic plane.  Standing against this evidence are two difficulties.  First, we must understand the serious discrepancy between the small timing age of the pulsar and its long kinematic timescale.  \\psr\\ now appears to be a prototypical ``young'' pulsar, with bright radio and X-ray emission from its associated pulsar wind nebula. If we assume that the pulsar is really $\\sim150$~kyr old, as required if the association is correct, then we must conclude that it spent the first 90\\% or more of its life looking much ``older'' and less energetic before being ``rejuvenated'' by some substantial increase in its spin-down rate $\\sim10^5$ years ago. While there seems no way to rule out the possibility that \\psr\\ was born with a weak magnetic field that has grown only recently, this appears to us to require unnatural fine-tuning.  We note that the magnetic field of \\psr\\ inferred from its spin-down rate, $4\\times10^{12}$~G, is not unusual for a young pulsar: it is within 10\\% of the magnetic field inferred for the 950~yr old Crab pulsar.  Second, we must understand how the remnant could have retained its circular symmetry over $1.5\\times10^5$~yrs while expanding two to three times faster to the south/south-east and decelerating to a velocity of a few hundred km/s or less.  Again, this appears to require unnatural fine-tuning of the parameters of the medium into which the remnant is expanding, with an extreme gradient in one direction and extreme uniformity in the transverse direction.  In summary, the unified model originally designed to explain the angular proximity of \\Gff, \\Gft, and \\psr\\ and the morphology of \\Gff\\ has failed both straightforward predictive tests initially proposed, and can be saved only by positing unusual evolution of {\\it both} the pulsar {\\it and} the remnant.  If instead we reject the hypothesis of a physical association, we find a more straightforward path. We conclude that \\psr\\ was most likely born about 15~kyr ago in or near the Crab-like remnant \\Gft---which it is escaping at $\\sim 5$~mas/yr (120~km/s at 5~kpc) and which it continues to power.  No new or unusual pulsar or remnant physics is required, nor any unusual properties for the local interstellar medium.  \\Gff\\ is an unrelated foreground or background remnant, whose not implausible angular association with \\Gft\\ in the dense Galactic center region of the sky is made more visually spectacular by its edge brightening toward the Galactic plane.  The Duck is no more a coherent body than is its northern cousin, Cygnus the Swan."
    },
    "0206/astro-ph0206202_arXiv.txt": {
        "abstract": "We present Hubble Space Telescope ultraviolet and optical STIS spectroscopy of the [WCL] planetary nebula central star \\cpd , obtained during its latest lightcurve minimum. The UV spectrum shows the central star's continuum light distribution to be split into two bright peaks separated by 0.10~arcsec. We interpret this finding as due to an edge-on disk or torus structure that obscures direct light from the star, which is seen primarily via its light scattered from the disk's rims or lobes. \\cpd\\ is an archetype of dual dust chemistry [WCL] planetary nebulae, which exhibit strong infrared emission features from both carbon-rich and oxygen-rich materials, and for which the presence of a disk harboring the O-rich grains had been suggested. Our direct observation of an edge-on occulting dust structure around \\cpd\\ provides strong support for such a model and for binary interactions being responsible for the correlation between the dual dust chemistry phenomenon in planetary nebulae and the presence of a hydrogen-deficient [WCL] Wolf-Rayet central star. ",
        "introduction": "WC Wolf-Rayet central stars of planetary nebulae ([WC] CSs of PNe) are H-deficient stars that exhibit strong ionic emission lines of helium, carbon and oxygen from their dense stellar winds. Amongst the coolest central stars in this group are \\cpd\\ (the nucleus of the PN He~3-1333) and the CS of He~2-113 (both classified as [WC10], Crowther et al. 1998). Cohen et al. (1986) found the mid-infrared KAO spectra of both these objects to show very strong unidentified infrared bands (UIBs -- usually attributed to polycyclic aromatic hydrocarbons, PAHs). Indeed, both the nebular C/O ratios (De Marco, Barlow \\& Storey 1997) and the ratio of UIB luminosity to total IR dust luminosity (Cohen et al. 1989) for these two objects are among the largest known. It was therefore a major surprise when mid- and far-IR Infrared Space Observatory spectra of these two objects showed the presence of many emission features longwards of 20~$\\mu$m that could be attributed to crystalline silicate and water ice particles (Barlow 1997, Waters et al.  1998a, Cohen et al. 1999), indicating a dual dust chemistry, i.e.  the simultaneous presence of both C-rich dust and O-rich dust. The dual dust chemistry phenomenon in PNe appears to show a strong correlation with the presence of a late WC ([WCL]) nucleus -- four out of six [WC8-11] nuclei studied by Cohen et al. (2002) showed similar dual dust chemistries. In the context of a single star scenario, this would point to a recent transition (within the last $\\sim$1000~yr) between the O- and the C-rich surface chemistries. However, the probability of finding a post-AGB object that had recently changed from an O-rich to a C-rich surface chemistry due to a third dredge-up event should be very low indeed. An alternative scenario envisages these systems as binaries (Waters et al. 1998a, Cohen et al. 1999, 2002), in which the O-rich silicates are trapped in a disk as a result of a past mass transfer event, with the C-rich particles being more widely distributed in the nebula as a result of recent ejections of C-rich material by the nucleus. Because of the quasi-periodic light variations shown by \\cpd , Cohen et al. (2002) suspected the presence of a precessing disk around it.  In this Letter we present the first direct evidence for an edge-on disk or torus around \\cpd, as revealed by recent HST STIS spectroscopy. ",
        "conclusions": "Our working hypothesis is that the observed spatial splitting of \\cpd 's continuum (Fig.~\\ref{fig:cpd_hen_ds9} and \\ref{fig:cuts}) is due to obscuration of the central object by a dusty disk or torus seen close to edge-on. Under this interpretation, the two continuum peaks separated by 0.10~arcsec are attributed to scattered light emerging from the upper and lower rims of the disk/torus, or from two lobes residing above and below the disk/torus, resembling the edge-on structure observed in the Red Rectangle, HD~44179 (Osterbart, Langer \\& Weigelt 1997) or the edge-on silhouette disk Orion 114-426 (McCaughrean et al. 1998). It seems possible that the STIS slit was oriented nearly normal to the PA of the disk, in agreement with the slit PA being only $\\sim$50~deg from the inferred axis of symmetry of the nebula. Our May 2001 STIS spectrum was obtained during a light minimum of \\cpd , the latest of three observed visual light declines having an apparent periodicity of $\\sim$5~yr (Cohen et al. 2002). Cohen et al. suggested that these light declines could be due to the precession of a nearly edge-on obscuring disk, or alternatively that they could be due to occultations of the central object by particularly dense clumps orbiting within the disk, in which case the apparent periodicity might be accidental.  One might argue that \\cpd's light variations could be due to random dust ejections, similar to those of R Coronae Borealis (R~CrB) stars. Although \\cpd 's apparent 5-year period can be confirmed only by more observations, we note that R~CrB light declines are much deeper ($\\sim$8~mag) and always associated with spectral variability (Clayton 1996). Even V348 Sgr (a [WC]-like CS not dissimilar to \\cpd ; Leuenhagen, Heber \\& Jeffery 1994) has deep declines and spectral variability, similarities to R~CrB characteristics which are not shared by \\cpd . Finally, resolved occulting structures, of the type found here around \\cpd , have never been detected around R~CrB stars.   As an alternative to disk precession being responsible for \\cpd 's visible light variability, the orbital motion of the star itself (clearly we have to assume the presence of a binary companion) might bring it in and out of alignment with a denser region in the disk. A similar behavior has been inferred by Van Winckel et al. (1999) for the binary post-AGB objects HD~52961 and HR~4049, for both of which circumstellar extinction variations were found to occur on the same timescales as the respective binary orbital periods of 1310~days and 429~days, though we note that their optical extinction variations show amplitudes of only 0.16--0.20~magnitudes, versus 1.6~magnitudes for \\cpd . If we assume the alleged companion star to have a similar mass to \\cpd\\ (and that both have masses of 0.6~M$_\\odot$), the orbital semi-major axis, for a circular orbit, would be 3.1~AU. The 0.10~arcsec angular separation between the two main scattered light components in the MAMA spectral image, are assumed to correspond to the top and bottom edges of a disk or torus with a projected disk semi-thickness of 67~AU, (for a distance of 1.35~kpc; De Marco et al. 1997) or to two lobes residing at 67~AU above and below the disk/torus. Cohen et al. (1999) estimated the 65-90~K oxygen-rich grains around \\cpd\\ to lie 1000~AU from the star, but this was based on an optically thin assumption. Our present results indicate that these cool O-rich grains are likely to be located much closer to the star, in the outer regions of an optically thick edge-on reprocessing disk. As noted above, a similar structure to that of \\cpd\\ is observed around the HD~44179 nucleus of the Red Rectangle, a non H-deficient post-AGB object that was found by Waters et al. (1998b) to exhibit a dual dust chemistry. Inspection of Figs.~2 and 3 of Cohen et al. (2002) shows the long wavelength peak of \\cpd 's infrared energy distribution to lie at the same wavelength as those of the three other [WCL] objects plotted, but also that this peak is much weaker relative to the 7.7~$\\mu$m UIB-like feature than is the case for the other three objects. If the crystalline silicate and long-wavelength continuum emission originates from an orbiting disk-like region around each object then this effect could be due to \\cpd 's disk being viewed edge-on, thereby significantly reducing the amount of radiation received from its disk.  One remaining issue to be considered is the very strong association between the dual dust chemistry phenomenon in PNe and the presence of a H-deficient central star, and what that might tell us about a common cause for the two phenomena. We might expect that some [WC] stars are single and simply evolve to be a [WC] CS as a result of the well-timed thermal pulses described by Herwig (2001), in which case they would not be expected to exhibit dual dust chemistries. Indeed, several [WCL] PNe show no evidence for dual dust chemistries (Cohen et al. 2002, Hony et al. 2001) and may well have originated via the well-timed thermal pulse mechanism. However, the current evidence points to binarity playing a key role in the origin of at least some H-deficient [WCL] nuclei {\\em and} the dual dust chemistry phenomenon, with mass transfers first leading to the creation of a massive circum-binary disk and later stripping most or all of the H-rich surface layer from the CS progenitor. Under this scenario, the B9/A0~III central star of the Red Rectangle might be interpreted as possessing only a very thin surface hydrogen layer, which will quickly be stripped away to yield a H-deficient [WCL] nucleus once the stellar effective temperature increases to high enough values ($>$ 20,000~K) for radiation pressure driven mass loss to become significant.  To conclude, our STIS observations indicate the presence of an edge-on occulting dust structure in the \\cpd\\ system -- further observations will be needed to determine how its appearance changes as the star returns to maximum. It would also be desirable for the orbital characteristics of the inferred \\cpd\\ binary system to be determined directly via radial velocity measurements of the [WC10] star's emission lines. Similar radial velocity measurements should also be obtained for the central stars of other dual dust chemistry [WCL] PNe, even if, as is likely in the case of He~2-113, the possible dust disks are not oriented edge-on to us."
    },
    "0206/astro-ph0206458_arXiv.txt": {
        "abstract": "{A finite element method for solving the resonance line transfer problem in moving media is presented. The algorithm works in three spatial dimensions on unstructured grids which are adaptively refined by means of an a posteriori error indicator. Frequency discretization is implemented via a first-order Euler scheme. We discuss the resulting matrix structure for coherent isotropic scattering and complete redistribution. The solution is performed using an iterative procedure, where monochromatic radiative transfer problems are successively solved. The present implementation is applicable for arbitrary model configurations with an optical depth up to $10^{3-4}$.  Results of Ly$\\alpha$ line transfer calculations for a spherically symmetric model, a disk-like configuration, and a halo containing three source regions are discussed. We find the characteristic double-peaked Ly$\\alpha$ line profile for all models with an optical depth $\\ga 1$. In general, the blue peak of the profile is enhanced for models with infall motion and the red peak for models with outflow motion. Both velocity fields produce a triangular shape in the two-dimensional Ly$\\alpha$ spectra, whereas rotation creates a shear pattern. Frequency-resolved Ly$\\alpha$ images may help to find the number and position of multiple Ly$\\alpha$ sources located in a single halo. A qualitative comparison with observations of extended Ly$\\alpha$ halos associated with high redshift galaxies shows that even models with lower hydrogen column densities than required from profile fitting yield results which reproduce many features in the observed line profiles and two-dimensional spectra. ",
        "introduction": "Hydrogen Ly$\\alpha$ as a prominent emission line of high redshift galaxies is important for the understanding of galaxy formation and evolution in the early universe. Recent narrow-band imaging and spectroscopic surveys used the Ly$\\alpha$ line to identify galaxies at very high redshift (e.g. Hu et al. \\cite{hu:etal98}; Kudritzki et al. \\cite{kudritzki:etal2000}; Rhoads et al. \\cite{rhoads:etal2000}; Fynbo et al. \\cite{fynbo:etal2000}, \\cite{fynbo:etal2001}).  But beside from being a good redshift indicator, Ly$\\alpha$ emission also bears information on the distribution and kinematics of the interstellar gas as well as the nature of the energy source. Ly$\\alpha$ observations of high redshift radio galaxies (e.g. Hippelein \\& Meisenheimer \\cite{hippelein:meisenheimer93}; van Ojik et al. \\cite{vanojik:etal96}, \\cite{vanojik:etal97}; Villar-Mart{\\'{\\i}}n et al. \\cite{villar-martin:etal99}) reveal extended Ly$\\alpha$ halos with sizes $>100$~kpc which are aligned with the radio jet. The line profiles are often double-peaked and the two-dimensional spectra point to complex kinematics involving velocities $>1000$~km$\\;\\mbox{s}^{-1}$.  The interpretation of Ly$\\alpha$ observations is difficult, because high-redshift radio galaxies tend to be in the center of proto clusters, where the radio jet interacts with a clumpy environment influenced by merging processes (e.g. Bicknell et al. \\cite{bicknell:etal2000}; Kurk et al. \\cite{kurk:etal2001}). Actually, the three-dimensional structure of the objects and the fact that Ly$\\alpha$ is a resonance line require detailed radiative transfer modeling.  The transfer of resonance line photons is profoundly determined by scattering in space and frequency. Analytical (Neufeld \\cite{neufeld90}) as well as early numerical methods (Adams \\cite{adams72}; Hummer \\& Kunasz \\cite{hummer:kunasz80}) were restricted to one-dimensional, static media. Only recently, codes based on the Monte Carlo method were developed which are capable to investigate the more general case of a multi-dimensional medium (Ahn et al. \\cite{ahn:etal2001}, \\cite{ahn:etal2002}).  In this paper, we introduce a finite element method for the solution of the resonance line transfer problem in three-dimensional, moving media and present the results of some simple, slightly optically thick model configurations. The basic, monochromatic code was originally developed by Kanschat (\\cite{kanschat96}) and is described in Richling et al. (\\cite{richling:etal2001}, hereafter Paper~I). The three-dimensional method is particularly useful for scattering dominated problems in inhomogeneous media. Steep gradients are resolved by means of an adaptively refined grid. Here, we only specify the extension from the monochromatic to the polychromatic problem including the implementation of the Doppler-effect and complete redistribution.  In Sect.~2, we review the equations for resonance line transfer in moving media. In Sect.~3, we describe some details regarding frequency discretization and the form of the resulting matrices for coherent scattering and complete redistribution and give a short outline of the complete finite element algorithm. Then, the results of a spherically symmetric model (Sect.~4), a disk-like configuration (Sect.~5) and a model with three separate source regions (Sect.~6) are presented. A summary is given in Sect.~7. ",
        "conclusions": "We presented a finite element code for solving the resonance line transfer problem in moving media. Non-relativistic velocity fields and complete redistribution are considered. The code is applicable to any three-dimensional model configuration with optical depths up to $10^{3-4}$.  We showed applications to the hydrogen Ly$\\alpha$ line of slightly optically thick model configurations ($\\tau\\le 10^2$) and discussed the resulting line profiles, Ly$\\alpha$ images and two-dimensional spectra. The systematic approach from very simple to more complex models gave the following results: \\begin{itemize} \\item An optical depth of $\\tau\\ga 1$ leads to the characteristic double peaked line profile with a central absorption trough as expected from analytical studies (e.g. Neufeld \\cite{neufeld90}). This form of the profile is the result of scattering in space and frequency. Photons escape via the line wings where the optical depth is much lower. \\item Global velocity fields destroy the symmetry of the line profile. Generally, the blue peak of the profile is enhanced for models with infall motion and the red peak for models with outflow motion. But there are certain velocity fields (e.g.~with steep gradients) and spatial distributions of the extinction coefficient, where the formation of a prominent peak is suppressed. \\item Double-peaked line-profiles show up as two emission regions in the two-dimensional spectra. Global infall or outflow leads to an overall triangular shape of the emission. Rotation produces a shear pattern resulting in banana-shaped emission regions for optical depths $\\ga 10$. \\item For non-symmetrical model configurations, the optical depth varies with the line of sight. Thus, the total flux, the depth of the absorption trough and the pattern in the two-dimensional spectra strongly depend on the viewing direction. \\end{itemize}  The applications demonstrate the capacity of the finite element code and show that the three-dimensional structure and the kinematics of the model configurations are very important. Thus, beside exploring higher optical depths up to the limits of our code, we will consider clumpy density distributions as well as dust absorption and try to model the Ly$\\alpha$ emission of individual high redshift galaxies. In addition, we intend to implement a second order method for the frequency discretization. Furthermore, we plan to overcome the difficulties in solving the line transfer problem for configurations with $\\tau>10^4$.  Therefore, we will extend our method and use the diffusion approximation in the most optically thick regions. In the other regions, the full frequency-dependent line transfer problem must be solved."
    },
    "0206/astro-ph0206491_arXiv.txt": {
        "abstract": "We report the discovery by the \\rxte\\ PCA of a second transient accreting millisecond pulsar, \\xtej, during regular monitoring observations of the galactic bulge region.  The pulsar has a spin frequency of 435 Hz, making it one of the fastest pulsars.  The pulsations contain the signature of orbital Doppler modulation, which implies an orbital period of 42 minutes, the shortest orbital period of any known radio or X-ray millisecond pulsar.  The mass function, $f_x = (1.278 \\pm 0.003)\\times 10^{-6} M_\\sun$, yields a minimum mass for the companion of between 0.013 and 0.017 $M_\\sun$, depending on the mass of the neutron star.  No eclipses were detected.  A previous X-ray outburst in June, 1998, was discovered in archival All-Sky Monitor data.  Assuming mass transfer in this binary system is driven by gravitational radiation, we constrain the orbital inclination to be in the range 30\\arcdeg--85\\arcdeg, and the companion mass to be 0.013--0.035 $M_\\sun$.  The companion is most likely a heated helium dwarf.  We also present results from the \\chandra\\ HRC-S observations which provide the best known position of \\xtej. ",
        "introduction": "Accreting neutron stars in low mass X-ray binaries (LMXBs) are conventionally thought to be the progenitors of millisecond or ``recycled'' radio pulsars \\citep{alpar82}.  Firm evidence supporting this theory remained elusive until the launch of NASA's {\\it Rossi X-ray Timing Explorer\\/} (\\rxte) in December, 1995.  The discovery of 300--600 Hz nearly coherent oscillations during thermonuclear X-ray bursts \\citep[e.g.,][]{stroh97} was a first solid indicator that neutron stars in LMXBs rotate rapidly.  This was followed by the discovery in April, 1998, of the first accreting millisecond pulsar, \\saxj, with a spin period of 2.5 ms and orbital period of 2.1 hr \\citep{wijnands-vdk98,chakmorgan98}.  This discovery convincingly established a link between accreting neutron stars and recycled pulsars.  The presence of quasi-periodic oscillations in the range 300--1300 Hz in many LMXBs has also been used to infer rapid neutron star spin \\citep[e.g.,][]{vdk00-khz}.  Binary evolution models are becoming more sophisticated \\citep{podsiad02}, but still involve significant assumptions about mass transfer and the effects of magnetic fields.  \\citet{kulkarni88} have questioned whether the birthrate of LMXBs can account for the number of millisecond radio pulsars.  On the other hand, there have been speculations that there should be a significant number of low-luminosity transient LMXBs in the galaxy \\citep{heise99,king00}, whose mass transfer and binary separation are driven primarily by the emission of gravitational radiation.  Although the discovery of \\saxj\\ provided convincing evidence that recycled pulsars can form in LMXBs, it is difficult to draw inferences on binary and stellar evolution based on a single case.  In this paper we report the discovery of an accreting millisecond pulsar, \\xtej, which was discovered by \\rxte\\ in regular monitoring of the galactic center region.  This is the fastest known accreting pulsar and the second of its kind to be found.  Recently, a third pulsar XTE J0929$-$314, was discovered \\citep{remillard02,galloway02-psr}.  Interestingly, all three systems have very low mass companions.  In \\S\\ref{sec:xte} and \\S\\ref{sec:chandra}, we present the discovery by \\rxte, and results of a short \\chandra\\ observation to determine the source position.  In \\S\\ref{sec:timing} we develop a pulsar timing solution, and in \\S\\ref{sec:spectral} we present basic spectral results. \\S\\ref{sec:disc} contains a discussion of the binary system properties.  In this paper we focus on the pulsar timing properties, and defer more detailed analyses of other issues to future work. ",
        "conclusions": "}   The binary orbital parameters allow us to estimate the properties of the companion.  The 42 minute orbital period immediately reveals that \\xtej\\ is a highly compact binary system.  \\xtej\\ has the shortest orbital period of all known millisecond pulsars (both X-ray and radio).  The amplitude of the modulations also determines the mass function of the pulsar shown in Table~\\ref{tab:timing}, defined as $f_x = (M_c \\sin i)^3 / (M_x + M_c)^2$, where $i$ is the binary inclination to our line of sight, and $M_x$ and $M_c$ are the masses of the pulsar and companion respectively.  Given $M_x$, the mass function provides the minimum possible value of $M_c$, which would occur when viewing the binary edge-on ($i = 90^\\circ$).  The minimum mass range shown in Table~\\ref{tab:timing} reflects a reasonable range of neutron star masses between 1.4--2.0 $M_\\sun$.  It is clear that the companion of \\xtej\\ is in the regime of very low mass dwarfs, of order 15 Jupiter masses.  This is a factor of $\\sim$3 smaller than the companion of \\saxj, which \\citet{bildchak00} speculate is a $0.05 M_\\sun$ brown dwarf.  It is reasonable to assume that the companion must fill its Roche lobe in order to transfer mass to the neutron star \\citep{eggleton83}. Combining the mass function and Roche lobe constraints results in a curve in the $M_c$ vs. $R_c$ plane, shown in Figure~\\ref{fig:mvr}. This line can be compared to the equations of state of other types of bodies, including hydrogen main sequence stars, brown dwarfs \\citep{chabrier00}, \\saxj\\ \\citep{bildchak00}, and a cold helium dwarf \\citep{zapolsky69,rappaport84}.  While none of the models intersect the trace of \\xtej, the hydrogen models are clearly unlikely.  The oldest brown dwarf model is unlikely, given the probability that irradiation by the compact object will cause bloating \\citep{bildchak00}.  A warm helium dwarf model, which would lie above the ``cold helium dwarf'' curve in Figure~\\ref{fig:mvr}, appears to be the most likely scenario, but there appear to be no calculations of such a configuration in the literature.  There were no X-ray eclipses or dips detected in the PCA light curves. We therefore put an upper limit of $i<85^\\circ$.  There is also no evidence of X-ray modulations at the binary period (at a limit of $~\\sim 0.5\\%$), which might have implied propagation through a scattering atmosphere in a near edge-on geometry \\citep{bildchak00}. On the other hand, a very low inclination is also unlikely, since that would imply a high companion mass, and thus a large mass transfer rate.  Following the reasoning of \\citet{bildchak00}, we find that the time-averaged mass accretion rate, based on the measured X-ray fluence and a recurrence time of 3.8 yr, is $\\dot{M}_x = 2.1\\times 10^{-11} M_\\sun$ yr$^{-1} d_{10}^2 m_{1.4}^{-1} T_{3.8}^{-1}$, where $d_{10}$ is the distance in units of 10 kpc, $m_{1.4} = M_x / (1.4 M_\\sun)$, and assuming a neutron star radius of 10 km.  On the other hand, the mass transfer in these systems is thought to be driven by gravitational radiation \\citep{king00,rappaport84}, in which case the mass transfer rate should be $\\dot{M}_{GR} = 1.2\\times 10^{-11} M_\\sun$ yr$^{-1} m_c^2 m_{1.4}^{2/3}$, where $m_c = M_c / (0.0137 M_\\sun$), and we have assumed the companion is nearly degenerate.  If the two mass transfer rates are equal, we arrive at the constraint, $\\sin i = 0.74 m_{1.4}^{3/2} d_{10}^{-1}$.  For distances within 15 kpc, this constraint implies inclinations of 30\\arcdeg--85\\arcdeg and companion masses in the range 0.013--0.035 $M_\\sun$, for neutron star masses between 1.4 $M_\\sun$ and 2.0 $M_\\sun$.  An interesting result of this constraint is that the distance is at least 7 kpc, significantly farther than \\saxj\\ \\citep[$d = 2.5$ kpc;][]{intzand01-1808}.  The peak luminosity of the persistent emission would then be $\\gtrsim 2\\times 10^{37}$ erg s$^{-1}$, an order of magnitude higher than that of \\saxj\\ \\citep{cui98}.  It is probable that \\xtej\\ is near the galactic center, in which case more pulsars like it should be detectable by the PCA bulge monitoring program over the next few years."
    },
    "0206/astro-ph0206172_arXiv.txt": {
        "abstract": "{We have determined mass loss rates and gas expansion velocities for a sample of 69 M-type irregular (IRV; 22 objects) and semiregular (SRV; 47 objects) AGB-variables using a radiative transfer code to model their circumstellar CO radio line emission.  We believe that this sample is representative for the mass losing stars of this type.  The (molecular hydrogen) mass loss rate distribution has a median value of 2.0$\\times$10$^{-7}$\\,M$_{\\odot}$\\,yr$^{-1}$, and a minimum of 2.0$\\times$10$^{-8}$\\,M$_{\\odot}$\\,yr$^{-1}$ and a maximum of 8$\\times$10$^{-7}$\\,M$_{\\odot}$\\,yr$^{-1}$.  M-type IRVs and SRVs with a mass loss rate in excess of 5$\\times$10$^{-7}$\\,M$_{\\odot}$\\,yr$^{-1}$ must be very rare, and among these mass losing stars the number of sources with mass loss rates below a few 10$^{-8}$\\,M$_{\\odot}$\\,yr$^{-1}$ must be small.  We find no significant difference between the IRVs and the SRVs in terms of their mass loss characteristics.  Among the SRVs the mass loss rate shows no dependence on the period.  Likewise the mass loss rate shows no correlation with the stellar temperature.  The gas expansion velocity distribution has a median of 7.0\\,km\\,s$^{-1}$, and a minimum of 2.2\\,km\\,s$^{-1}$ and a maximum of 14.4\\,km\\,s$^{-1}$.  No doubt, these objects sample the low gas expansion velocity end of AGB winds.  The fraction of objects with low gas expansion velocities is very high, about 30\\% have velocities lower than 5\\,km\\,s$^{-1}$, and there are objects with velocities lower than 3\\,km\\,s$^{-1}$: \\object{V584~Aql}, \\object{T~Ari}, \\object{BI~Car}, \\object{RX~Lac}, and \\object{L$^2$~Pup}.  The mass loss rate and the gas expansion velocity correlate well, a result in line with theoretical predictions for an optically thin, dust-driven wind.  In general, the model produces line profiles which acceptably fit the observed ones.  An exceptional case is \\object{R~Dor}, where the high-quality, observed line profiles are essentially flat-topped, while the model ones are sharply double-peaked.  The sample contains four sources with distinctly double-component CO line profiles, i.e., a narrow feature centered on a broader feature: \\object{EP~Aqr}, \\object{RV~Boo}, \\object{X~Her}, and \\object{SV~Psc}. We have modelled the two components separately for each star and derive mass loss rates and gas expansion velocities.  We have compared the results of this M-star sample with a similar C-star sample analysed in the same way.  The mass loss rate characteristics are very similar for the two samples.  On the contrary, the gas expansion velocity distributions are clearly different.  In particular, the number of low-velocity sources is much higher in the M-star sample.  We found no example of the sharply double-peaked CO line profile, which is evidence of a large, detached CO-shell, among the M-stars.  About 10\\% of the C-stars show this phenomenon. ",
        "introduction": "It has been firmly established that mass loss from the surface is a very important process during the final stellar evolution of low- and intermediate-mass stars, i.e., on the asymptotic giant branch (AGB). The mass loss seems to occur irrespective of the chemistry (C/O$<$1 or $>$1) or the variability pattern (irregular, semi-regular, or regular) of the star.  Beyond these general conclusions the situation becomes more uncertain \\citep{olof99}.  Of importance for comparison with mass loss models and for the understanding of AGB-stars in a broader context (e.g., their contribution to the chemical evolution of galaxies) is to establish the mass loss rate dependence on stellar parameters, such as main sequence mass, luminosity, effective temperature, pulsational pattern, metallicity, etc., and its evolution with time for individual sources.  A crude picture has emerged where the average mass loss rate increases as the star evolves along the AGB, and where the final mass loss rate reached increases with the main sequence mass.  In addition, there is evidence of time variable mass loss \\citep{haleetal97,maurhugg00}, and even highly episodic mass loss \\citep{olofetal00}.  There is a dependence on luminosity in the expected way, i.e., an increase with increasing luminosity, but it is uncertain how strong it is.  The same applies to the effective temperature where a decrease with increasing temperature is expected. Regular pulsators clearly have higher mass loss rates than stars with less regular pulsation patterns.  The dependence of the total mass loss rate on metallicity appears to be weak, but the dust mass loss rate decreases with decreasing metallicity.  See \\citet{habi96} for a summary of evidence in favour of this general outline.  The mechanism behind the mass loss remains unknown, even though there are strong arguments in favour of a wind which is basically pulsation-driven, and where the highest mass loss rates and gas expansion velocities are reached through the addition of radiation pressure on dust \\citep{hoefdorf97,wintetal00b}.  A way to study this problem is to use samples of low mass loss rate stars for which stellar parameters can be reasonably estimated using traditional methods.  These samples also contain objects with quite varying pulsational characteristics, and has, as it turned out, quite varying circumstellar characteristics.  \\citet{olofetal93a} presented such a study of low mass loss rate C-stars using CO multi-transition radio data.  These data were subsequently analysed in more detail by \\citet{schoolof01} using a radiative transfer model.  In the same spirit \\citet{kersolof99} presented a major survey of CO radio line emission from irregularly variable (IRV) and semiregularly variable (SRV) M-type AGB-stars.  They increased the number of IRVs (22 detections), in particular, and SRVs (43 detections) detected in circumstellar CO emission substantially ($\\approx$60\\% of the SRVs and all but one of the IRVs were detected for the first time). \\citet{youn95} and \\citet{groeetal99} have made extensive surveys of short-period M-Miras.  In this paper we use the radiative transfer method of \\citet{schoolof01} to estimate reasonably accurate mass loss rates and gas expansion velocities for the \\citet{kersolof99} sample.  Comparisons between these properties and other stellar properties are done, as well as comparisons with the results for the C-star sample. ",
        "conclusions": "We have determined mass loss rates and gas expansion velocities for a sample of 69 M-type IRVs (22 objects) and SRVs (47 objects) on the AGB using a radiative transfer code to model their circumstellar CO radio line emission.  We believe that this sample is representative for the mass losing stars of this type.  The uncertainties in the estimated mass loss rates are rather low within the adopted stellar/circumstellar model, typically less than $\\pm$50\\%.  However, a sensitivity analysis shows that for these low mass loss rate stars there is a considerably uncertainty due to the stellar luminosity, the size of the CO envelope, the CO abundance, and as usual the distance to the source. We find that the mass loss rates determined by the detailed radiative transfer analysis differ by almost an order of magnitude from those obtained by published mass loss rate formulae.  The (molecular hydrogen) mass loss rate distribution has a median value of 2.0$\\times$10$^{-7}$\\,M$_{\\odot}$\\,yr$^{-1}$, and a minimum of 2.0$\\times$10$^{-8}$\\,M$_{\\odot}$\\,yr$^{-1}$ and a maximum of 8$\\times$10$^{-7}$\\,M$_{\\odot}$\\,yr$^{-1}$.  M-type IRVs and SRVs with a mass loss rate in excess of 5$\\times$10$^{-7}$\\,M$_{\\odot}$\\,yr$^{-1}$ must be very rare, and in this respect the regularity and amplitude of the pulsation plays an important role. We also find that among these mass losing stars the number of sources with mass loss rates below a few 10$^{-8}$\\,M$_{\\odot}$\\,yr$^{-1}$ must be small.  We find no significant difference between the IRVs and the SRVs in terms of their mass loss characteristics.  Among the SRVs the mass loss rate shows no dependence on the period.  Thus, for these non-regular, low-amplitude pulsators it appears that the pulsational pattern plays no role for the mass loss efficiency.  We have determined temperatures for our sample stars by fitting blackbody curves to their spectral energy distributions.  These blackbody temperatures have been shown to correlate reasonably well with the stellar effective temperatures.  The mass loss rates of our stars show no correlation at all with these stellar blackbody temperatures.  The gas expansion velocity distribution has a median of 7.0\\,km\\,s$^{-1}$, and a minimum of 2.2\\,km\\,s$^{-1}$ and a maximum of 14.4\\,km\\,s$^{-1}$.  No doubt, these objects sample the low gas expansion velocity end of AGB winds.  The fraction of objects with low gas expansion velocities is high, about 30\\% have velocities lower than 5\\,km\\,s$^{-1}$.  There are four objects with gas expansion velocities lower than 3\\,km\\,s$^{-1}$: \\object{V584~Aql}, \\object{T~Ari}, \\object{BI~Car}, \\object{RX~Lac}, and \\object{L$^2$~Pup}. These objects certainly deserve further study.  We find that the mass loss rate and the gas expansion velocity correlate well, $\\dot{M}$$\\propto$$v_{\\rm e}^{1.4}$, even though for a given velocity (which is well determined) the mass loss rate may take on a value within a range of a factor of five (the uncertainty in the mass loss rate estimate is lower than this within the adopted circumstellar model). The result is in line with theoretical predictions for an optically thin, dust-driven wind.  A more detailed test of the CO modelling is provided by the shape of the line profiles.  In general, the fits are acceptable, but there is a trend that the model profiles, in particular the $J$=1$\\rightarrow$0 ones, are more flat-topped, or even weakly double-peaked, than the observed ones.  An exceptional case is \\object{R~Dor}, where the high-quality, observed line profiles are essentially flat-topped, while the model ones are sharply double-peaked.  Acceptable fits are obtained  by increasing the distance to the star or by artificially decreasing the size of the CO envelope.  The sample contains four sources with distinctly double-component CO line profiles: \\object{EP~Aqr}, \\object{RV~Boo}, \\object{X~Her}, and \\object{SV~Psc} (all SRVs).  We have modelled the two components separately for each star and derive mass loss rates and gas expansion velocities using the same circumstellar model as for the rest of the sample.  The resulting mass loss rates and gas expansion velocities show the same positive correlation as that of the other objects.  At present, the exact nature(s) of these objects is unknown.  We have compared the results of this M-star sample with a similar C-star sample.  The mass loss rate distributions are comparable, suggesting no dependence on chemistry for these types of objects. Likewise, the mass loss rates of the C-stars show no correlation with stellar temperature or period. The gas expansion velocity distributions though are clearly different.  The fraction of low velocity sources is much higher in the M-star sample.  In both cases there is a correlation between mass loss rate and gas expansion velocity, although the detailed relations are different.  Our crude estimates of the dust properties, through the gas-grain collision heating term, indicate that the two samples have similar gas-to-dust ratios and that these differ significantly from that of high mass loss rate C-stars.  This also means that the gas-CSEs due to low mass loss rates are cooler than expected from a simple extrapolation of the results for \\object{IRC+10216}.  Finally, we find no example of the sharply double-peaked CO line profile, which is evidence of a large, detached CO-shell, among the M-stars.  About 10\\% of the C-stars show this phenomenon."
    },
    "0206/astro-ph0206344_arXiv.txt": {
        "abstract": "Using Gemini QuickStart infrared observations of the central $22''$ of M33, we analyze the stellar populations in this controversial region. Based on the slope of the giant branch we estimate the mean metallicity to be $-0.26 \\pm 0.27$, and from the luminosities of the most luminous stars, we estimate that there were two bursts of star formation $\\sim 2$ and $\\sim 0.5$ Gyr ago.  We show that the stellar luminosity function not only has a different bright end cutoff, but also a significantly different slope than that of the Galactic bulge, and suggest that this difference is due to the young stellar component in M33.  We combine our infrared Gemini data with optical HST-WFPC2 measurements revealing a CMD populated with young, intermediate, and old age stellar populations. Using surface brightness profiles from $0.1''$ to $18'$, we perform simple decompositions and show that the data are best fit by a three-component, core + bulge + disk model.  Finally, we find no evidence for radial variations of the stellar populations in the inner $3-10''$ of M33 based on a spatial analysis of the color-magnitude diagrams and luminosity functions. ",
        "introduction": "\\label{sec:introduction}  With the advent of space-based telescopes, such as the Hubble Space Telescope (HST), and large aperture ground-based telescopes with adaptive optics (AO), such as Gemini and VLT, the number of galaxies beyond the Milky Way (MW) and its dwarf companions for which detailed studies can be made is gradually increasing.  Of the nearby galaxies, M33 (NGC 598, the Triangulum Nebula) is one of the best for studying the stellar content of spiral galaxies.  Outside of the MW, it is one of the closest and brightest spirals visible, surpassed only by M31.  M31 is more luminous and slightly closer, but it's higher inclination angle of 77\\degr\\ (compared to 56\\degr\\ for M33) makes it more difficult to separate the contributions from different stellar populations.  However, is M33 the prototype or the exception when it comes to late-type spiral galaxies?  It contains a healthy population of halo globular clusters \\citep{Schommer1991}, yet it is not certain whether there is a matching bulge.  This simple point is crucial for understanding the role of a bulge in galaxy formation, and its relationship to the halo and globular clusters.  As of 1991, the status of M33's bulge was ``controversial'' according to a review of literature by \\citet{vandenbergh1991}.  There have since been many papers on M33, but the number of authors who find a bulge seems balanced by an equal number who do not. \\citet{Bothun1992} argued against a traditional bulge, based on the lack of a power-law contribution to his 12\\micron\\ surface brightness measurements. Although his $B$-band observations do show an excess of light inside $3'$, he finds no satisfactory $r^{1/4}$ fit.  Based on infrared (IR) observations of the central $2.5' \\times 8'$ which show a clustering of stars around the nucleus, \\citet{Minniti1993, Minniti1994} take the opposite side.  They find a de Vaucouleurs profile fits all but the inner $1'$ of their surface photometry.  \\citet{Regan1994} are also in favor of a bulge in M33.  Their infrared surface brightness measurements of the central $15' \\times 30'$ seem reasonably well fit by an exponential disk plus $r^{1/4}$ profile.  In contrast, \\citet{McLean1996} cite an unchanging IR luminosity function from $45''$ to $1.5'$ as evidence against a significant bulge population beyond $45''$.  Thus, a decade after van den Bergh's review, the controversy remains unresolved.  In this paper we use IR Gemini-North observations with QUIRC/Hokupa'a to study the stellar populations in the inner regions of M33.  In section \\ref{sec:observations} we describe our observations and observing procedure.  Section \\ref{sec:data_reduction} details the data reduction, photometric techniques, and calibration.  We give the azimuthally averaged surface brightness profile in Section \\ref{sec:decomposition} and, after combining our observations with those of \\citet{Regan1994}, present simple two- and three- component model decompositions.  We describe artificial star tests performed to understand completeness and observational effects as a function of a star's position and luminosity in Section \\ref{sec:artstartests}.  We use the color-magnitude-diagram (CMD) to estimate the metallicities and ages of the stars in our field in Section \\ref{sec:cmds}.  We present the stellar luminosity function (LF) in Section \\ref{sec:lfs}, and compare it to the LF measured in the Galactic bulge.  In section \\ref{sec:comparison} we compare our observations to the recent work in the infrared by \\citet{Davidge2000a} with the CFHT AO system, and combine our data with optical HST-WFPC2 measurements made by \\citet{Mighell2002}. We look for radial variations in the stellar properties in Section \\ref{sec:radial_variations}. In Section \\ref{sec:blending} we perform a theoretical analysis of blending on our own and previous observations following the procedures of \\citet{Renzini1998}.  Finally we give a summary of our conclusions in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have used Gemini-North to study the stellar populations in the central regions of M33.  The surface brightness profiles from $0.1''$ to $18'$, formed from the combination of our data and those of \\citet{Regan1994}, show that the data need to be modeled using a three-component, core + spheroid + disk model.  The best-fit parameters are listed in Table \\ref{tab:decomposition}.  These high-resolution observations allow us to accurately measure individual stars to $K \\sim 21$.  Artificial star tests (\\S \\ref{sec:artstartests}) show that our completeness is relatively uniform across the field (50\\% at $K=21$), although within $3''$ from the nucleus the completeness is dramatically lower due to severe crowding.  Artificial fields are used to understand the observational effects associated with adaptive optics measurements in crowded fields.  Based on the slope of the giant branch in the infrared color magnitude diagram, we estimate the mean metallicity to be $-0.26 \\pm 0.27$ (\\S \\ref{sec:metallicity}).  Using the bolometric luminosities and density of stars on the AGB, we hypothesize two bursts of star formation; at $\\sim 0.5$ and $\\sim 2$ Gyr ago (\\S \\ref{sec:age}).  We note however, that this component of young stars may have influenced our metallicity estimate due to the sensitivity of the GB slope on age.  The stellar luminosity function in M33 is shown to be significantly different from that measured in the Galactic Bulge as viewed through Baade's Window (\\S \\ref{sec:lfs}).  The difference in their maximum luminosities is due to differences in ages of the two regions.  We speculate that this is also the origin of their different slopes as well.  In section \\ref{sec:comparison} we compare our data with previous observations.  Recent work by \\citet{Davidge2000a} is in good agreement with our data, although our observations go $>2$ magnitudes deeper (\\S \\ref{sec:davidge}).  We also combine our data with optical HST-WFPC2 measurements \\citep{Mighell2002}, and present the optical-IR CMD in \\S \\ref{sec:mighell}.  This CMD clearly shows that the central regions of M33 are composed of young, intermediate, and old aged stellar populations.  Dividing the inner $\\sim 8.5''$ ($\\sim 35$ pc) into equal area rings around the nucleus, we look for radial variations in the stellar properties (\\S \\ref{sec:radial_variations}).  However, based on the distribution of stellar luminosities and the morphology of the CMD, we find that all stars are consistent with being drawn from a single population.  In the last section (\\S \\ref{sec:blending}) we perform calculations to estimate the severity of blending at various imaging resolutions and locations in M33.  Using the formulations of \\citet{Renzini1998} and our composite surface brightness profile, these calculations call into question some previous claims of very luminous stars in the central regions of M33."
    },
    "0206/astro-ph0206422_arXiv.txt": {
        "abstract": "We present an initial analysis of a new \\xmm\\ observation of NGC 1399, the central elliptical galaxy of the Fornax group. Spectral fitting of the spatially resolved spectral data of the EPIC MOS and pn CCDs reveals that a two-temperature model (2T) of the hot gas is favored over single-phase and cooling flow models within the central $\\sim 20$~kpc. The preference for the 2T model applies whether or not the data are deprojected. The cooler component has a temperature ($\\sim 0.9$~keV) similar to the kinetic temperature of the stars while the hotter component has a temperature ($\\sim 1.5$~keV) characteristic of the virial temperature of a $\\sim 10^{13}\\msun$ halo. The two-phase model (and other multitemperature models) removes the ``Fe Bias'' within $r\\la 20$~kpc and gives $\\fe/\\solar\\approx 1.5-2$. At larger radii the iron abundance decreases until $\\fe/\\solar\\sim 0.5$ for $r\\sim 50$~kpc. The Si abundance is super-solar (1.2-1.7 solar) within the central regions while $\\si/\\fe\\approx 0.8$ over the entire region studied. These Fe and Si abundances imply that $\\approx 80\\%$ of the Fe mass within $r\\sim 50$~kpc originates from Type Ia supernovae (SNIa).  This SNIa fraction is similar to that inferred for the Sun and therefore suggests a stellar initial mass function similar to the Milky Way. ",
        "introduction": "\\label{intro}  NGC 1399, the central galaxy of the Fornax group, is one of the brightest elliptical galaxies in X-rays and has been the subject of numerous X-ray studies. Although previous \\rosat\\ and \\asca\\ studies demonstrated that the hot gas within $r\\sim 30$~kpc is not isothermal \\citep[e.g.,][]{rang95,jone97,buot99a}, the data could not distinguish between single-phase and multiphase models. Moreover, we showed that an ``Fe Bias'' for systems like NGC 1399 occurs if an isothermal gas is assumed when in fact the spectrum consists of multiple temperature components with values near 1 keV \\citep[e.g.,][]{buot00a,buot00c}. We argued that this bias is primarily responsible for the very sub-solar Fe abundances found in the central regions of the (X-ray) brightest ellipticals and groups.  The combined spatial and spectral resolution of the \\xmm\\ CCDs allows for unprecedented mapping of the temperatures and elemental abundances of the hot gas in galaxies, groups, and clusters. These capabilities complement the high spatial resolution ($\\sim 1\\arcsec$) of \\chandra\\ which has already provided interesting constraints on the emission from discrete sources near the center of NGC 1399 \\citep{ange01}. The higher energy resolution, sensitivity, and larger field-of-view of the \\xmm\\ EPIC CCDs are better suited for constraining the spatial and spectral properties of the diffuse hot gas out to radii well past the optical extent of the galaxy (i.e., out to $r\\sim 50$~kpc assuming a distance of 21~Mpc using the results of \\citet{tonr01} for $H_0=70$~\\kmsmpc).  We present initial results for the temperature and metal abundances of the hot gas obtained from a new \\xmm\\ observation of NGC 1399. Detailed discussions of the data reduction, spectral fitting, and results for the gravitating mass will appear in \\citet{buot02d} and \\citet{lewi02b}. ",
        "conclusions": "\\label{conc}  We have performed a spatially resolved spectral analysis of the \\xmm\\ EPIC data of NGC 1399 which suggests the presence of two dominant phases of hot gas within the central $\\sim 20$~kpc. This two-phase description is consistent with 2T models we fitted previously to the \\asca\\ data  \\citep{buot99a}. However, the \\asca\\ data lacked the spatial resolution of \\xmm\\ and could not distinguish between a radially varying single-phase medium, a two-phase medium, or a cooling flow within $r\\sim 30$~kpc. These results for NGC 1399 are supported by previous analyses of its X-ray surface brightness profile which argues for separate galaxy and group components \\citep[e.g.,][]{ikeb96,paol02b} similar to other centrally E-dominated groups \\citep{mulc98}.  Since the temperature structure of NGC 1399 obtained by \\asca\\ is characteristic of other central E-dominated galaxy groups \\citep{buot99a,buot00a,buot99b,alle00}, the two-phase structure revealed by \\xmm\\ for NGC 1399 likely applies to these other systems. Moreover, since the centers of several clusters show similar multitemperature structure with \\asca\\ \\citep[e.g.,][]{buot99b,alle01} and similar multi-component surface brightness structure \\citep[see][and references therein]{etto00}, we might expect two-phase structures in the centers of these systems. A recent \\xmm\\ analysis of M87 is consistent with this suggestion \\citep{mole01b}.  The two-phase model provides a physical explanation for the gas density and temperature profiles seen in ``cooling flows''. The hotter component is associated with the ambient group or cluster gas while the cooler component is associated with the stellar ejecta of the dominant central galaxy. Gasdynamical models of groups and clusters constructed by \\citet{brig99b,brig02a} demonstrate that relatively cool ejecta from stellar mass loss combined with the hotter ambient gas can reproduce observed single-phase temperature profiles in groups and clusters. But our results suggest that these two phases are not fully mixed.  Understanding the details of the interaction of these phases may shed light on the mystery of cooling flows. Using the results of \\citet{math89}, the rate of energy input by stellar ejecta is, $\\approx\\alpha_{\\star}M_{\\star}\\sigma^2$, where $\\alpha_{\\star}\\approx 5\\times 10^{-20}$~s$^{-1}$ is the specific stellar mass-loss rate. If within $r\\sim 5$~kpc we take $M_{\\star}\\approx 2\\times 10^{11}\\msun$ for the stellar mass and $\\sigma=310$ km s$^{-1}$ \\citep{jone97}, we infer the rate of energy input into the central $\\sim 5$~kpc of NGC 1399 from stellar ejecta to be $\\approx 2\\times 10^{40}$~\\ergs. This crude estimate is only a factor of two less than the X-ray luminosity of the hot component (and a factor of four less than the cool component). If the energy input from stellar ejecta does not completely suppress mass drop-out, it may at least delay the episodes of AGN feedback. Problems with such heating models are discussed by \\citet{brig02b}.   The two-phase model (as well as other multitemperature models) removes the ``Fe Bias'' within the central regions ($r\\la 20$~kpc) and gives $\\fe/\\solar\\approx 1.5-2$. These super-solar values exceed the stellar values in typical elliptical galaxies \\citep[e.g.,][]{trag00a} and therefore allow for significant enrichment of the hot gas from Type Ia supernovae (SNIa) \\citep[cf.][]{arim97}. For example, using the results of \\citet{gibs97} and our measured value of $\\si/\\fe \\approx 0.8$ solar, we infer that SNIa have contributed $\\approx 80\\%$ of the iron mass within $r\\sim 50$~Kpc in NGC 1399. This SNIa fraction is similar to that inferred for the Sun and therefore suggests a stellar initial mass function similar to the Milky Way as advocated by Renzini and others \\citep[e.g.,][]{renz93,renz97,wyse97}. We close by noting that although correcting for the ``Fe Bias'' partially removes the ``Iron Discrepancy'' noted by \\citet{arim97}, chemical models of elliptical galaxies without cooling flows predict central iron abundances even larger than we have measured for NGC 1399 \\citep[e.g.,][]{brig99a}."
    },
    "0206/astro-ph0206087_arXiv.txt": {
        "abstract": "We report the discovery of very high latitude molecular gas in the edge-on spiral galaxy, NGC~5775. Emission from both the J=1-0 and 2-1 lines of $^{12}$CO is detected up to 4.8~kpc away from the mid-plane of the galaxy. NGC~5775 is known to host a number of HI supershells. The association of the molecular gas (M$_{H_2,F2}$ = 3.1$\\times$10$^7$~M$_{\\sun}$) reported here with one of the HI supershells (labeled F2) is clear, which suggests that molecular gas may have survived the process which originally formed the supershell. Alternatively, part of the gas could have been formed in situ at high latitude from shock-compression of pre-existing HI gas. The CO J=2-1/J=1-0 line ratio of 0.34$\\pm$40\\% is significantly lower than unity, which suggests that the gas is excited subthermally, with gas density a few $\\times$ 10$^2$~cm$^{-3}$. The molecular gas is likely in the form of cloudlets which are confined by magnetic and cosmic rays pressure. The potential energy of the gas at high latitude is found to be 2$\\times$ 10$^{56}$~ergs and the total (HI + H$_2$) kinetic energy is 9$\\times$ 10$^{53}$~ergs. Based on the energetics of the supershell, we suggest that most of the energy in the supershell is in the form of potential energy and that the supershell is on the verge of falling and returning the gas to the disk of the galaxy. ",
        "introduction": "Galactic HI shells and supershells, distinguished by whether their initial energy requirement is less or more than 10$^{53}$~ergs, were first studied by \\citet{hei79,hei84}. The large energies found in the supershells ($>$ 10$^{53}$~ergs) imply that these structures must have a tremendous influence on the structure of the interstellar medium. In addition, supershells that break through the gaseous disk to reach high galactic latitudes may be a source of star formation in the halo. For example, supershells may act as ''chimneys'' through which hot gas from the disk funnels to the halo \\citep[e.g., the ''chimney model'',][]{nor89}. This hot gas may cool and eventually form stars in the halo.  Alternatively, the molecular gas in the supershell may reach high-latitude and directly provide raw material for star formation in the halo. Although which of these two scenarios is at work cannot be distinguished easily, the study of molecular gas in supershells still provides an important clue to high-latitude star formation. Molecular gas at high latitude also presents an important aspect in the understanding of the global evolution of the interstellar medium in spiral galaxies. In the Milky Way Galaxy, the study of supershells is hindered by difficulties with distance determination and the resulting confusion. In external galaxies, these problems are minimized.  NGC~5775 is an edge-on ($i$ = 86\\arcdeg), infrared-bright (L$_{FIR}$ = 2.6$\\times$10$^{10}$~L$_{\\sun}$) galaxy at a distance of 24.8~Mpc. \\citet{irw94} observed this galaxy and its face-on neighbour NGC~5774 in HI using the VLA and provided models for their HI distributions. She showed that the two galaxies may be engaging in an early phase of interaction, with two HI bridges connecting them. Numerous HI arcs and extensions beyond the disk of NGC~5775 are also observed. Six HI supershells were cataloged by \\citet{lee98}. In a multi-wavelength study of NGC~5775, \\citet{lee01} report spatial correlations of HI, radio continuum, X-rays and far-infrared emission at the positions of the three largest HI supershells, labeled F1 through F3 in Fig.~\\ref{hi}.  \\begin{figure} \\plotone{figure1.ps} \\caption{HI column density contours superimposed on the Digital Sky Survey image of the vicinity of NGC~5775. The white crosses indicate the centres of NGC~5775 (left centre), NGC~5774 (upper right) and IC~1070 (bottom). Contour levels are at 1, 5, 10, 17.5, 25, 40, 60, and 82.6$\\times$10$^{20}$~cm$^{-2}$. The HI supershells are labeled F1 to F3. \\label{hi}} \\end{figure} ",
        "conclusions": "This paper reports the detection of high-latitude molecular gas in NGC~5775. The shell-like distribution of the CO emission coincides exactly with that of the HI supershells, suggesting that we have detected the molecular shell associated with the HI supershell. The existence of the molecular shell means that molecules are not destroyed during the ejection of the supershell and is entrained in the expanding flow to reach high latitude. Some of the molecular gas may have been formed in situ, via shock-compression of pre-existing HI gas. The CO J=2-1/1-0 line ratio (0.34$\\pm$40\\%) suggests that the gas density in the supershell is low and the gas is subthermally excited. The molecular gas is probably in the form of cloudlets which are confined by magnetic and cosmic rays pressure. Based on energetics ground, we proposed that the supershell may be at a stage where it is about to plunge towards the disk of the galaxy, returning the gas to the bulk of the gas reservoir of the galaxy."
    },
    "0206/astro-ph0206278_arXiv.txt": {
        "abstract": "Intracluster magnetic fields with $ \\sim \\mu \\rm G $ strength induce Faraday rotation on the cosmic microwave background (CMB) polarization.  Measurements of this effect can potentially probe the detailed structure of intracluster magnetic fields across clusters, since the CMB polarization is a continuously varying field on the sky, in contrast to the conventional method restricted by the limited number of radio sources behind or inside a cluster.  We here construct a method for extracting information on magnetic fields from measurements of the effect, combined with possible observations of the Sunyaev-Zel'dovich effect and $X$-ray emission for the same cluster which are needed to reconstruct the electron density fields. Employing the high-resolution magneto-hydrodynamic simulations performed by Dolag, Bartelmann \\& Lesch (1999) as a realistic model of magnetized intracluster gas distribution, we demonstrate how our reconstruction technique can well reproduce the magnetic fields, i.e., the spherically averaged radial profiles of the field strength and the coherence length. ",
        "introduction": "The origin and evolution of cosmic magnetic fields are still unclear and outstanding problems.  Various observational techniques have consistently revealed that most clusters of galaxies are pervaded by magnetic fields of $B \\sim \\rm O(1) \\mu \\rm G$ strength (see e.g., Carilli \\& Taylor 2002 for a review).  Recently, Clarke, Kronberg \\& B\\\"ohringer (2001) have again drawn a robust conclusion that the intracluster hot plasma typically has a $B \\sim 5-10 ~\\mu$G field assuming a coherence length of $10 \\rm{kpc}$ from the Faraday rotation measurements of 16  low-z ($ z \\leq 0.1$) clusters selected to be free of unusual radio halos.  Since the rotation measure arises from the integration of the product of the electron density and the line-of-sight component of magnetic fields, it is generally difficult to extract information on the magnetic fields only from the rotation measure without introducing any assumptions on the gas distribution and the field configuration.  The first systematic study of the structure of magnetic fields over a single cluster was performed by Kim et al. (1990) based on the rotation measures of 18 sources close in angular position to the Coma cluster, giving the estimation of $ B \\sim 2\\mu$G with the coherence length of $10 \\rm{kpc}$ which is indicated from the magnetic field reversal scale.  In a subsequent study, Feretti et al. (1995) discovered smaller coherence lengths down to 1 kpc from the rotation measures of the extended radio galaxy near the Coma cluster center, whereby the field strength estimation was modified to $B \\sim 8\\mu \\rm G$ to explain the measured rotation angle. Thus, it is crucial for estimating the magnetic field strength robustly to determine the magnetic field coherence length which does not necessarily match the coherence length of the rotation measures. To study the detailed structure of the intracluster magnetic fields for any clusters, the high-resolution measurements of the Faraday rotation should at least  be performed. However, there are some limitations for the conventional methods  because of the lack of the number of radio sources behind or inside a cluster and possible contributions of the intrinsic Faraday rotation. Moreover, we should break the degeneracy of the rotation measure between the electron density and the magnetic field strength using some additional information for estimating the magnetic field strength.  Recently, Takada, Ohno \\& Sugiyama (2001) proposed that the magnetized intracluster gas similarly induces a Faraday rotation effect on the linearly polarized radiation of the cosmic microwave background (CMB) generated at the decoupling epoch of $z\\approx 1000$ (see e.g., Hu \\& White 1997 for a study of the primary CMB polarization). Hereafter, to distinguish the origin of the CMB anisotropies, we use ``primary'' to indicate the anisotropies from the recombination epoch while ``secondary'' to those from the cluster formation epoch. They also calculated the angular power spectra of this secondarily induced polarization under the simple assumption of a uniform field strength with $\\sim 0.1\\mu~{\\rm G}$ across a cluster and suggested that the measurements could be used to set constraints on the average properties of the intracluster magnetic fields.  Cooray, Melchiorri \\& Silk (2002) computed the secondary power spectrum including a circularly polarized contribution characterized by the Stokes-V parameter that is induced by possible relativistic plasma in clusters via the Faraday rotation. However, as a more interesting and realistic possibility, one may imagine that the secondary effect can be in principle used to map the detailed structure of the magnetic fields in an individual cluster, since the CMB polarization is a continuously varying field on the sky.  In this paper, therefore, we study a method for reconstructing the magnetic fields from measurements of the Faraday rotation effect on the CMB polarization, combined with accessible observations of the Sunyaev-Zel'dovich (SZ) effect and $X$-ray emission for the same cluster to reconstruct the electron density distribution. For this purpose, it is crucial to consider a realistic magnetic field configuration as well as a plausible gas distribution as suggested by the currently favored formation scenario of galaxy clusters. Hence, as for a model of the magnetized cluster, we employ here high-resolution magneto-hydrodynamic simulation results performed by Dolag, Bartelmann \\& Lesch (1999), whereby we can simulate maps of the CMB polarization including the Faraday rotation effect as well as maps of the SZ effect and the thermal $X$-ray emission. Using those simulated maps, we demonstrate how well our method can reconstruct the magnetic fields. In particular, we try to clarify the relationship between the coherence lengths of the magnetic fields and the rotation measure. The coherence length should give a new insight into the nature and evolutionary history of magnetic fields. For example, if the magnetic fields have the coherence scale as large as or larger than the cluster size, the seed fields should be generated at the early stage of the universe (see e.g., Grasso and Rubinstein 2001 for a review), while smaller coherence lengths may imply the seed fields originated from galaxies within the cluster (e.g., Kronberg 1996).  This paper is organized as follows. In \\S 2, we refer to the cluster models which are used to demonstrate the magnetic field reconstruction. In \\S 3, we briefly review the primary CMB polarization map generated at the decoupling epoch and the Faraday rotation effect. In \\S 4, assuming spherical symmetry of clusters, we show a method for reconstructing the density fields of the clusters from SZ effect and thermal $X$-ray emission which are directly calculated from simulated clusters. In \\S 5, we develop a way to reconstruct the coherence length and strength of the magnetic fields combined with the previously reconstructed electron density fields. Finally, \\S 6 is devoted to a summary and discussion. ",
        "conclusions": "In this paper we have constructed a new method for reconstructing magnetic fields in galaxy clusters from the Faraday rotation effect on the CMB polarization, combined with possible observed maps of the $X$-ray emission  and SZ effect on the CMB. Our results imply that the CMB polarization can be potentially used to reconstruct detailed radial profiles of the coherence length and strength  of the magnetic fields. It was shown that the coherence length estimated from the rotation measure matches that of the magnetic fields. Therefore, we do not need any other information than the rotation measure for estimating the magnetic field coherence length. This coherence length is not only an important quantity for determining the magnetic field strength but also could reveal the origin of the initial seed fields.  To reconstruct the field strength, we consider the dispersion of the rotation measure fields in the annulus of a given projected radius, motivated by the random walk process caused by random orientations of the magnetic fields for a cell with the coherence length. However, the deviation of the field strength from the spherically averaged value may also increase the dispersion of the rotation measure, which we do not consider in this paper. This increase of the dispersion should lead to the increase of the reconstructed magnetic field strength. Nevertheless, the field strengths are well reconstructed in the three cluster models, and are reconstructed  within a factor of a few even around the area where the dispersion of the field strength is large. Anyway, we expect that  the method constructed in this paper will be a powerful tool for probing the intracluster magnetic fields.  Finally, we comment on the feasibility of this method. It is a great challenge for current technology to detect the secondary effect of the intracluster magnetic fields on the CMB polarization. The rotation angle becomes about $1-10^\\circ$ at a frequency of $10$GHz (e.g. $\\Delta_{RM} \\sim 1-10 {}^{o} (10 \\rm {GHz}/ \\nu)^2$ ). The sensitivity of the detector, which is needed to detect the Faraday rotation,  is of order $1\\,\\mu K$. The angular resolution needed to reconstruct the structure of the magnetic fields is estimated from the minimum coherence length of the magnetic fields in the simulation cluster as $\\sim 50 \\rm{kpc} \\sim 20''$ (at z $\\sim 0.1$). The frequency dependence of the Faraday rotation can be also used to discriminate the effect from other secondary signals. Many observations are ongoing and planned for measuring the CMB polarization. We expect that future extensive observations of the CMB polarizations will allow reconstructions of intracluster magnetic fields with sufficient accuracy, which should give a crucial key to understanding the origin of intracluster magnetic fields."
    },
    "0206/astro-ph0206093_arXiv.txt": {
        "abstract": "{We use an analytic approach to study the susceptibility of the {\\sc Planck} Low Frequency Instrument radiometers to various systematic effects. We examine the effects of fluctuations in amplifier gain, in amplifier noise temperature and in the reference load temperature. We also study the effect of imperfect gain modulation, non-ideal matching of radiometer parameters, imperfect isolation in the two legs of the radiometer and back-end $1/f$ noise. We find that with proper gain modulation $1/f$ gain fluctuations are suppressed, leaving fluctuations in amplifier noise temperature as the main source of $1/f$ noise. We estimate that with a gain modulation factor within $\\pm 1\\%$ of its ideal value the overall $1/f$ knee frequency will be relatively small ($< 0.1 $Hz). ",
        "introduction": "Introduction}  {\\sc Planck}\\footnote{{\\sc Planck} homepage: http://astro.estec.esa.nl/Planck/} is a European Space Agency (ESA) satellite mission to map spatial anisotropy and polarization in the Cosmic Microwave Background (CMB) over a wide range of frequencies with an unprecedented combination of sensitivity, angular resolution, and sky coverage (Bersanelli et al. \\cite{bersanelli1}). Following the breakthrough of the COBE discovery of CMB anisotropy (Smoot et al. \\cite{smoot}, Bennett et al. \\cite{bennett}), and the MAP\\footnote{MAP homepage: http://map.gsfc.nasa.gov} satellite launched in June 2001, {\\sc Planck} will be the third generation space mission dedicated to CMB observations.  The data gathered by these missions will revolutionise modern cosmology by a precise determination of the fundamental cosmological parameters which govern the present expansion rate of the universe, the average density of the universe, the amount of dark matter, and the nature of the seed fluctuations from which all structures in the universe arose (see, e.g., Hu et al. \\cite{Hu}, Scott et al. \\cite{scott}, White et al. \\cite{white} for recent reviews on CMB anisotropy) and will provide at the same time full sky surveys at essentially unexplored frequencies, with fundamental implications for a large area of problems in astrophysics (De Zotti et al. \\cite{dezotti}).  {\\sc Planck} consists of a High Frequency Instrument (HFI) (Puget et al. \\cite{puget}) and a Low Frequency Instrument (LFI) (Mandolesi et al. \\cite{mandolesi}) observing the sky through a common telescope. While the MAP radiometers measure temperature differences between two widely separated regions of the sky through a pair of symmetric back-to-back telescopes, the {\\sc Planck}~LFI radiometers are designed to measure differences between the sky signal and a stable internal cryogenic reference load. The LFI scheme takes advantage of the presence in the focal plane of the HFI front end unit, which is cooled to approximately 4 K as an intermediate cryo stage for the 0.1 K bolometer detectors.  The LFI radiometer design is a modified correlation receiver (Blum \\cite{blum}, Colvin \\cite{colvin}, Bersanelli et al. \\cite{bersanelli2}), realised with High Electron Mobility Transistor (HEMT) amplifiers at 30, 44, 70 and 100 GHz. The modification is that the temperature of the reference load can be made significantly different from the sky temperature. To compensate for the offset (a few K in nominal conditions), a {\\sl gain modulation factor}, $r$, is used to null the output signal in order to minimise sensitivity to RF gain fluctuations and achieve the lowest white and $1/f$ noise in the output.  Obtaining data streams with low $1/f$ noise is of primary importance in order to achieve the LFI scientific objectives. In fact excessive $1/f$ noise would degrade the quality of the measured data (Janssen et al. \\cite{janssen}) by increasing the effective rms noise and the uncertainty in the power spectrum at low multipole values. Such effects can be avoided if the post detection knee frequency $f_k$ (i.e. the frequency at which the $1/f$ contribution and the ideal white noise contribution are equal) is significantly lower than the spacecraft rotation frequency ($f_{\\rm spin}\\sim 0.017$ Hz). For values of $f_k$ greater than $f_{\\rm spin}$ it is possible to mitigate such effects by applying appropriate {\\em destriping} and {\\em map making} algorithms\\footnote{see Burigana et al. \\cite{burigana1}, Delabrouille \\cite{delabrouille}, Maino et al. \\cite{maino2}, for details about {\\em destriping} and Dor\\'e et al. \\cite{dore}, Natoli et al. \\cite{natoli}, for details about {\\em map-making} algorithms.} to the time ordered data (Maino et al. \\cite{maino1}).  If the knee frequency is sufficiently low (i.e. $f_k\\leq 0.1$~Hz), with the application of such algorithms it is possible to maintain the increase in rms noise within few \\% of the white noise, and the power increase at low multipole values (i.e. $l\\leq 200$) at a very low level ($\\sim$ two order of magnitude less than the CMB power). If, on the other hand, the knee frequency is high (i.e. $\\gg 0.1$~Hz) then even after destriping the degradation of the final sensitivity is of several tenths of \\% and the excess power at low multipole values is significant (up to the same order of the CMB power for $f_k\\sim 10$~Hz, Bersanelli et al., \\cite{bersanelli2002}). Therefore, careful attention to instrument design, analysis, and test is essential in order to achieve a low $1/f$ noise knee frequency.  In this paper, we analyse the most important systematic effects due to non-ideal behaviour of components in the LFI radiometer signal chain, and estimate their impact on the post detection knee frequency. In section~\\ref{sec:analytic_model_LFI} we present a general analytical description of the {\\sc Planck}-LFI radiometers to derive formulas for the radiometer power output and sensitivity in the two cases of perfectly balanced and slightly unbalanced radiometer. Here we also show that under quite general assumptions, the radiometer sensitivity does not depend on the reference load temperature. In section~\\ref{sec:systematic_effects} we analyse the impact of various systematic effects on the post-detection knee frequency showing that with proper gain modulation it is possible to keep the radiometer $1/f$ noise to a very low level also in the presence of different non-ideal behaviours (e.g., gain and noise temperature imbalance, imperfect gain modulation). For sake of conciseness, we transfer part of the formalism to the appendices. Finally, in section~\\ref{sec:conclusions} we summarise our results and discuss briefly their implications for Planck observations. ",
        "conclusions": "}  In this paper we have discussed the pseudo-correlation architecture adopted for the radiometers of the {\\sc Planck}-LFI instrument and we have studied the sensitivity of the measured signal to various systematic effects. In our treatment we have considered both the ideal case of a perfectly balanced radiometer and the effect of small mismatches in the various radiometer parameters.  The first result is that the radiometer sensitivity does not depend on the level of the reference load temperature; even in the case of a slight imbalance in the radiometer parameters the dependence on $T_y$ is at the level of $\\partial \\Delta T /\\partial T_y\\sim 10^{-5}$ which is negligible. The only mismatch which has a first-order impact on $\\Delta T$ is the noise temperature mismatch of the front-end amplifiers; our analysis shows that it is possible to maintain the sensitivity degradation below 1\\% with a noise temperature match better than 5\\%.  With proper gain modulation ($r=r^*_0$) the $1/f$ noise in the radiometer output is determined mainly by noise temperature fluctuations in the front-end amplifiers, with a knee frequency of few mHz, provided that the front-end amplifier amplitude match is better that $\\sim\\pm 0.5$~dB. Such a high level of $1/f$ noise suppression depends on the gain modulation factor, which must be determined with an accuracy better than $\\pm 0.2\\%$.  If the accuracy on $r$ is less than $\\pm 1\\%$ then gain fluctuations become the major source of $1/f$ noise; if $\\epsilon_r$ is kept in the range $\\pm 1\\%$ we expect values of the knee frequency of the order of 50~mHz which can be easily handled by destriping algorithms. The presence of a small amount of leakage in the first hybrid does not significantly modify this conclusion; with the correct choice of $r$, one can make the sensitivity to gain fluctuations vanish and null the average output of the radiometer simultaneously.  The effect of gain fluctuations in the back-end amplifiers can be made negligible by the fast front-end switching between sky and reference load signals. The LFI baseline of 4096~Hz for the phase switch frequency has been chosen to guarantee a high suppression level of the $1/f$ noise from back-end amplifiers.  In general our analysis demonstrates the effectiveness of the gain modulation concept applied to this form of radiometer. The estimate of the knee frequency given by equation (\\ref{eq:kneefreq}) is quite low and relatively immune to small imperfections in radiometer balance. The modified correlation radiometer scheme reduces the knee frequency by more than two orders of magnitude, compared to a total power radiometer of similar bandwidth and intrinsic transistor fluctuations. For such small residual knee frequency (of $\\sim 0.1$~Hz) it will be possible to remove efficiently the effects during data analysis.  We have also studied the sensitivity of the radimeter to changes in the reference load temperature, $T_y$. Our analysis provides a framework in which thermal stability requirements on the LFI reference loads can be evaluated.  A refinement of the present analysis for the determination of $f_k$ will be pursued in the future by software simulations of the radiometer functions to accurately study the combined effect of all components. Finally, laboratory measurements of a prototype radiometer working under conditions close to those of Planck mission constitute the most important checks for the ultimately understanding of the behaviour of {\\sc Planck}~LFI radiometers regarding the $1/f$ type noise and possible further effects. Results of preliminary laboratory measurements performed of {\\sc Planck}-LFI prototype radiometers will be presented in forthcoming publications."
    },
    "0206/astro-ph0206436_arXiv.txt": {
        "abstract": "Radio astronomy has provided evidence for the presence of ionized atmospheres around almost all classes of non-degenerate stars. Magnetically confined coronae dominate in the cool half of the Hertzsprung-Russell diagram. Their radio emission is predominantly of non-thermal origin and has been identified as gyrosynchrotron radiation from mildly relativistic electrons, apart from some coherent emission mechanisms. Ionized winds are found in hot stars and in red giants. They are detected through their thermal, optically thick radiation, but synchrotron emission has been found in many systems as well. The latter is emitted presumably by shock-accelerated electrons in weak magnetic fields in the outer wind regions. Radio emission is also frequently detected in pre-main sequence stars and protostars, and has recently been discovered in brown dwarfs. This review summarizes the radio view of the atmospheres of non-degenerate stars, focusing on energy release physics in cool coronal stars, wind phenomenology in hot stars and cool giants, and emission observed from young and forming stars. ",
        "introduction": "Stellar radio astronomy has matured over the past two decades, driven in particular by discoveries made with the largest and most sensitive radio interferometers.  Radio emission is of great diagnostic value as it contains telltale signatures not available from any other wavelength regime. Some of the detected radio emission represents the highest-energy particle populations (MeV electrons) yet accessible on stars, the shortest (sub-second) detectable time scales of variability and energy release, and probably refers most closely to the primary energy release responsible for coronal heating. This review is to a large extent devoted to demonstrating the ubiquity of high-energy processes in stars as revealed by radio diagnostics.  Stellar radio sources include thermal and non-thermal magnetic coronae, transition regions and chromospheres, stars shedding winds, colliding-wind binaries, pre-main sequence stars with disks and radio jets, and embedded young objects visible almost exclusively by their radio and millimeter-wave emission. Most recent additions to the zoo of objects are brown dwarfs, and with the increasingly blurred transition from  stars to brown dwarfs to giant planets like Jupiter and Saturn, even the magnetospheres of the latter may have to be considered a manifestation of magnetic activity in the widest sense. To keep the discussion somewhat focused, this paper concentrates on physical processes in magnetic coronae, but includes,  in a more cursory way, atmospheres of young and forming stars and winds of hot stars. We do not address in detail the large phenomenology of extended and outflow-related sources such as radio jets, HII regions, masers, Herbig-Haro objects,  and the diverse millimeter/submillimeter phenomenology, e.g., molecular outflows and dust disks. Compact stellar objects (white dwarfs, neutron stars, black holes) are not considered here.  Inevitably - and fortunately - much of the knowledge gained in stellar astronomy is anchored in solar experience. The privilege of having a fine specimen - and even an exemplary prototype - next door is unique among various fields of extrasolar astrophysics, being shared since recently only by the related field of extrasolar planetary astronomy. Detailed solar studies, even in-situ measurements of the solar wind, are setting high standards for studies of stellar atmospheres, with a high potential reward. Solar radio astronomy has been reviewed extensively in the literature. For detailed presentations, we refer to \\citet{dulk85} and \\citet{bastian98}.  The maturity of stellar radio astronomy is demonstrated by a number of  review articles on various subjects; a non-exhaustive list for further reference includes \\citet{andre96, bastian90, bastian96, bookbinder88, bookbinder91, dulk85, dulk87, guedel94a,  hjellming88, kuijpers85a, kuijpers89a, lang90, lang94, lestrade97, linsky96, melrose87, mullan85, mullan89, mutel96, phillips91r, seaquistrev93, vdoord96a}, and \\citet{white96, white00}. A natural starting point for this review, even if not strictly adhered to, is Dulk's comprehensive 1985 Annual Review  article that summarizes the pre- and early-Very Large Array view of stellar (and solar) radio emission. Meanwhile, two dedicated conferences, the first one in Boulder in 1984 (proceedings edited by \\citealt{hjellming85}) and the second held in Barcelona in 1995 (proceedings edited by \\citealt{taylor96}), provided a rich forum to discuss new developments; together, they beautifully illustrate the progress made over the past decades. ",
        "conclusions": "Stellar radio astronomy has matured over the past few decades to a science that is  indispensable for our understanding of stellar atmospheres. Historical milestones include, among many others, the discovery of steady and flaring non-thermal and polarized emission in cool stars, testifying to the importance of highly energetic processes; the recognition that these phenomena are ubiquitous in many classes of convective-envelope stars; observations of very large, apparently stable magnetospheric structures, unlike anything known from the Sun, around various types of magnetically active stars such as T Tau stars, Bp stars, dMe stars, or RS CVn binaries; the discovery of non-thermal emission produced in (wind-collision) shocks of hot-star atmospheres; gyromagnetic and flaring emission from deeply embedded protostellar objects, testifying to the importance of magnetic fields back to the earliest moments of a stellar life; and flaring radio emission from sub-stellar objects not previously thought to support stellar-like convective outer envelopes. Radio methodology has become a standard to estimate magnetic fields in cool stars, to determine mass loss in stars with ionized winds, to spatially resolve and map structures at the milliarcsecond level, and to simply prove the presence of magnetic fields through polarization measurements.  Far  from being an auxiliary science to research at other wavelengths, stellar radio astronomy should prepare to address outstanding problems to which it has unique access, although more sensitive instruments are needed. Questions of particular interest include: Are there relevant high-energy processes and magnetic fields in class 0 protostars? Are accretion processes important for the high-energy mechanisms and the generation of large-scale magnetic fields? Are there magnetic fields in hot stars, and what role do they play in the  winds? Are brown dwarfs usually quiescent radio emitters? Do they maintain stable magnetic fields? What is the structure of their coronae? Are there intra-binary magnetic fields in close binary stars? How do large magnetospheres couple to the more compact X-ray coronae? Are quiescent coronae fed by numerous (micro-)flares? \\vskip 0.5truecm  Acknowledgements: It is a pleasure to thank Marc Audard, Arnold Benz, Stephen Skinner, Kester Smith, and Stephen White for helpful discussions, and  Philippe Andr\\'e, Tim Bastian, Arnold Benz,  Maria Contreras, Jeremy Lim, Robert Mutel, Tom Ray, Luis Rodr\\'{\\i}guez, and Stephen White for providing figure material. The introductory Einstein quotation is cited from \"Einstein sagt\", ed A Calaprice and A Ehlers (1997, Munich: Piper) and from \"The Expanded Quotable Einstein\", ed A Calaprice (2000, Princeton: Princeton University Press)."
    },
    "0206/astro-ph0206399_arXiv.txt": {
        "abstract": "{We present and analyze quiescent $UBVRI$ light curves of the classical symbiotic binary YY Her. We show that the secondary minimum, which is clearly visible only in the quiescent $VRI$ light curves, is due to ellipsoidal variability of the red giant component. Our simple light curve analysis, by fitting of the Fourier cosine series, resulted in a self-consistent phenomenological model of YY Her, in which the periodic changes can be described by a combination  of the ellipsoidal changes and a sinusoidal changes of the nebular continuum and line emission. ",
        "introduction": "Symbiotic stars are long-period interacting binaries made up of an evolved red giant and a hot companion surrounded by an ionized nebula. The hot component in the vast majority of systems seems to be a luminous ($\\sim 1000\\, \\mathrm{L_{\\sun}}$) and hot ($\\sim 10^5\\, \\mathrm{K}$) white dwarf powered by thermonuclear burning of the material accreted from its companion's wind. Depending on the accretion rate, these systems can be either in a steady burning configuration or undergo hydrogen shell flashes, which in many cases can last for decades and even centuries due to low mass of the white dwarf. The latter can be the case for non-eruptive symbiotics (such as RW Hya or V443 Her), in which the hot components maintain roughly constant luminosity and temperature for many decades (Miko{\\l}ajewska  \\cite{mik97}).  In many systems, however, the hot components show multiple outburst activity with timescales of a few/several years, which cannot be simply accounted for by the thermonuclear models. Among them, Z And is one of  the best studied (e.g. Miko{\\l}ajewska \\& Kenyon \\cite{mk96}, and references therein), and the symbiotic stars with multiple outburst activity are often referred to as Z And-type systems.  The photometric history of YY Her since 1890 has been studied in detail by Munari et al. (\\cite{munari97a}, hereafer M97a) who revealed, in addition to Z And-type multiple outburst activity, a periodic fluctuations with $P=590^\\mathrm{d}$ and a visual amplitude of $\\la 0.3^\\mathrm{m}$. The orbital nature of this periodicity has been confirmed by subsequent UBV photometric observations (Munari et al. \\cite{munari97b}; Tatarnikova et al. \\cite{tatar01}, hereafter M97b and T01, respectively). In particular, the photometric and spectroscopic study of YY Her during  the major outburst in 1993  and its decline (Tatarnikova et al. \\cite{tatar00}, hereafer T00; T01) has shown that the minima are due to obscuration of the hot component and the ionized nebula. The return to quiescence in 1997/98 was accompanied by significant changes in the shape of light curves, and in particular by the appearance of a secondary minumum in the $VRI$ light curves (T01; Hric et al. \\cite{hric01}, hereafter H01). T01 interpreted the secondary minimum in terms of ellipsoidal changes of the red giant, whereas H01 argued that the secondary minimum is caused by an eclipse of the red giant by a circumstellar envelope  around the hot component.  In this paper we present new $UBVRI$ photometric observations obtained during the secondary minimum in 2001, and demonstrate that the quiescent light curves of YY Her can be succesfully reproduced by a combination of ellipsoidal changes of the cool giant and sinusoidal changes of the nebular continuum and line emission. ",
        "conclusions": "Our very simplistic light curve analysis, namely independent fitting of the Fourier cosine series to the quiescent $UBVR'I'$ light curves, resulted in a self-consistent phenomenological model of YY Her, in which the periodic changes can be described by variations of the nebular emission along the orbit (the $\\cos \\phi$ term) combined with the ellipsoidal variability of the red giant (the $\\cos 2\\phi$ term).  YY Her is not the only symbiotic system with the secondary minima. Similar light curve behaviour is observed in CI Cyg, whose quiescent $UBVRI$ light curves (Miko{\\l}ajewska \\cite{mik01}) are almost identical with the light curves shown in Fig.~\\ref{lc2}. The secondary minima are also present in quiescent near infrared light curves of BF Cyg (Miko{\\l}ajewska et al. \\cite{mik02}).  We also think that the alternative interpretation of the secondary minimum in terms of the eclipse of the cool giant by the circumstellar envelope of the hot component (H01) is implausible. First, it cannot be the ionized nebular envelope because to cause eclipses it must be optically thick in the optical/red continuum whereas the presence of a prominent Balmer jump indicates that it is not. Second, to account for the amplitudes of the secondary minimum, it must screen $\\sim 20$\\,\\% of the red giant's projected surface. Such a big and hot screen, however, should dominate the quiescent optical continuum, which is not observed. Finally, we can also exclude eclipses of the cool giant by an accretion disc. A geometrically and optically thick disc seen nearly edge-on could mimic a cool source, and screen significant part of the giant's surface. Such a disc, however, would also completely hide the hot white dwarf wheras the quiescent shortwavelength {\\it IUE} spectrum clearly shows a hot continuum (M97b). Similarly, a hot rising continuum is present in quiescent shortwavelength {\\it IUE} spectra of CI Cyg and BF Cyg.  One of fundamental questions in relation of symbiotic binaries is the mechanism that powers the Z And-type multiple outburst activity observed in many classical symbiotic systems. The quiescent spectra of these systems can mostly be fitted by a hot stellar source (most likely a white dwarf powered by thermonuclear burning of the material accreted from its companion) and its ionizing effect on a nebula (e.g. M{\\\"u}rset et al. \\cite{murset91}; M97b, Miko{\\l}ajewska \\& Kenyon \\cite{mk96}). Their outburst activity, however, with time scales of a few/several years cannot be simply accounted by the thermonuclear models. A promising interpretation of this activity involves changes in mass transfer and/or accretion disc instabilities. Detection of an ellipsoidal hot continuum source during outbursts of CI Cyg, AX Per, YY Her, AS 338 and other Z And-type systems suggests the presence of an optically thick accretion disc (Miko{\\l}ajewska \\& Kenyon \\cite{mk92}; Esipov et al. \\cite{esipov00}; T00; Miko{\\l}ajewska et al. \\cite{mik02}) strongly supports this interpretation.  Related to this problem is the process of mass transfer -- Roche lobe overflow or stellar wind -- and the possibility of an accretion disc formation. The presence of secondary minima in the light curves of YY Her, apparently due to ellipsoidal changes of the red giant, as well as in a few other similar systems provides important clues. Although it is very premature to claim that all symbiotic systems with the Z And-type activity do have tidally distorted giants, and -- at least during active phase -- accretion discs, whereas the non-eruptive systems (such as RW Hya) do not, the former is certainly the case for YY Her, CI Cyg and a few other active systems.  Systematic searches for ellipsoidal variations in both active and non-eruptive symbiotic stars are necessary to address the problem, and confirm or exclude tidally distorted giants. The observations must be, however, carried in the red and near-IR range where the cool giant dominates the continuum light. We also note that any optical/red light curve analysis must be supported by spectroscopic information about contamination of the broad-band photometry by the nebular component."
    },
    "0206/astro-ph0206350_arXiv.txt": {
        "abstract": " ",
        "introduction": "Since the first observation of the Cosmic Microwave Background (CMB) by Penzias \\& Wilson (1965), many observations have been made to measure the CMB absolute temperature both at high and low frequencies. The most precise values are given by the FIRAS instrument aboard the COBE satellite (Mather et al. 1999).  In this work we present the complete data of the measurements of the CMB absolute temperature. We linger in particular on the most recent observations, and thus obtaining a critical database of the measures. We analyse the papers in the literature to recognize the main causes of error and to evaluate separately the magnitude of the systematical and statistical errors. As well known, the CMB spectrum provides significant informations on physical processes in the universe at very high redshifts (e.g. Danese \\& Burigana 1993 and references therein). The construction of a complete and manageable database is a first step for a versatile statistical analysis of CMB spectrum data and their comparison with the theoretical predictions for the distorted spectra.  \\medskip  For the following discussion, the whole set of the CMB spectrum observations has been divided in five sections corresponding to different frequencies ranges. Separately, we analyse the measures of the CMB temperature obtained from the molecular observations (section~\\ref{sez:CN}) and those by FIRAS and COBRA (section~\\ref{sez:satelliti}). This division in frequency ranges is mainly related to the different problems which dominate the determination of the CMB temperature at different frequencies:  \\begin{enumerate}  \\item $0.408<\\nu<1.0$ GHz (section \\ref{sez:0.4-1}), where the uncertainties due to the observation site choice and to the determination of the Galactic contribution dominate; \\item $1.4<\\nu<2.0$ GHz (section \\ref{sez:1-2}), where the uncertainty on the Galactic temperature dominates; \\item $2.3<\\nu<9.4$ GHz (section \\ref{sez:2-9}), where, whilst the Galactic contribution decreases, the uncertainty on the atmospheric temperature becomes important in the measure of $T_{CMB}$; \\item $10<\\nu<37$ GHz (section \\ref{sez:10-37}), where experiments both from ground and from balloons have been made. The last ones allow, observing at an altitude of about 25 km, to reduce the problem of the atmosphere contaminating the measuraments from the ground; \\item above 50 (section \\ref{sez:50-}). In this section only the values of the CMB temperature from ground and balloon are reported, without further analyses, since in this range the very accurate results from FIRAS are available. \\end{enumerate} ",
        "conclusions": "We have reviewed all the existing measures of the CMB absolute temperature. After a discussion of the main contamination relevant for the whole set of observed frequencies, we have considered in detail the different sources of contamination relevant at the different frequency bands for ground based, balloon and space observations. We have compared the relative weights of statistical and systematic uncertainties, by focussing in particular on the most recent observations. On the basis of this analysis, we have constructed a complete and reasoned database of CMB absolute temperatures, that allows to easily recognize the relevant informations on the different observations. The simple database format permits a reading of the data through any text editor or through programs for data handling. We have then implemented a set of tools in FORTRAN, that allow to easily select the desired set of measures and discriminate between the quoted statistical and systematic errors. This constitutes a the first step for a versatile comparison of the existing data with the theoretical predictions for the distorted spectra in order to derive constraints on physical processes at very high redshifts.  \\bigskip \\bigskip  \\noindent {\\bf Acknowledgements.} It is a pleasure to thank M.~Bersanelli, L.~Danese, G.~De~Zotti, N.~Mandolesi and G.~Palumbo for useful and stimulating discussions.  \\bigskip \\bigskip"
    },
    "0206/astro-ph0206166_arXiv.txt": {
        "abstract": "According to two estimated relations between the initial period and the dynamo-generated magnetic dipole field of pulsars, we calculate the statistical distributions of pulsar initial periods. It is found that proto-pulsars are very likely to have rotation periods between 20 and 30 ms, and that most of the pulsars rotate initially at a period $< 60$ ms. ",
        "introduction": "Pulsars provide us an unusual physical condition to increase our knowledge of the nature. As yet, one of the essential parameters, the initial periods $P_0$ of proto-pulsars,\\footnote{ Pulsars could be neutron stars or strange stars (e.g., Xu 2001). We discuss in this paper both possibilities of proto-pulsars being proto-neutron stars and proto-strange stars. }% which may reveal precious information of dynamical supernova process, are poorly known. Actually there are 3 efforts to estimate $P_0$ in the literatures. 1, The period $P_0$ can be found via $T=P/[(n-1)\\dot P][1-(P_0/P)^{n-1}]$ if the age $T$ and the braking index $n\\equiv {\\Omega \\ddot \\Omega/ \\dot{ \\Omega}^2}$ ($\\Omega=2\\pi/P$ the angular velocity of rotation) are measured (e.g., Kaspi et al. 1994), where $P$ is the rotation period observed. This method needs constant braking indices, which are likely to vary with time (Xu \\& Qiao 2001). 2, Monte Carlo simulation is used to study the pulsar ``current'' in the magnetic field - period diagram (e.g., Lorimer et al. 1993); the initial period of ``injected'' pulsars into the population are adjusted so as to sustain the number of pulsars observed at longer periods. 3, Initial periods can be inferred for pulsars that reside within composite supernova remnants which are powered by the pulsar spindown energy (van der Swaluw \\& Wu 2001). The derived $P_0$ via these 3 ways are 10-60 ms, $\\sim 300$ ms, and 37-82 ms (sometimes several hundred milliseconds), respectively. In addition, Lai et al. (2001) studied the implication of the apparent spin-kick alignment in the Crab and Vela pulsars, and found that the initial period should be $\\la 1$ ms in both the electromagnetic rocket model and the hydrodynamic natal model, or $<10$ s in the asymmetric neutrino emission model, in order to produce a high kick velocity.  In this paper, we try an alternative effort to derive the initial period through a relation between the initial period and the dynamo-originated magnetic field of pulsars. The magnetic field $B$ is assumed not to change in our discussion, since both observation and theory imply that a pulsar's $B$-field does not decay significantly during the rotation-powered phase (Bhattacharya et al. 1992 and, e.g., Xu \\& Busse 2001). Therefore the field strength observed for a normal pulsar represents its initial value. On another hand, the strong magnetic field is supposed to be created by dynamo action during the proto-pulsar stage (Thompson \\& Duncan 1993, Xu \\& Busse 2001). We may obtain the initial period if we can find an estimated relation between $B$ and $P_0$ during the dynamo episode. ",
        "conclusions": "The statistical distributions of pulsar initial periods are inferred based on the estimated relation between the initial periods and the magnetic dipole magnetic fields. It is found that proto-pulsars are very likely to have rotation periods between 20 and 30 ms, and that most of the pulsars rotate initially at a period $< 60$ ms. The pulsar initial periods derived in this way are comparable with the previous estimates by the first and the third methods, but are significantly smaller than that by the second method (section 1).  When looking at the pulsar distribution in the $P-\\dot P$ diagram, one generally concludes that the magnetic fields decays at a time scale of $10^{6-7}$ years since the number density of pulsars with low fields and old ages are larger than that of pulsars with high fields and young ages. However, calculations of ohmic decay of dipolar magnetic fields by Sang \\& Chanmugam (1987) suggest that neutron star magnetic fields may not decay significantly during the rotation-powered phases. The field of a strange star can also hardly decay because of the high magnetic Reynolds number and possible color superconductivity (e.g., Xu \\& Busse 2001). Observations indicate that the field decays only if the pulsar is in the accretion phase (e.g., Taam \\& van den Huevel 1986). If the field does not change significantly, how can we understand the apparent ``field decay'' in the $P-\\dot P$ diagram? We may explain this discrepancy by studying pulsar current in the $P-\\dot P$ diagram, according to eq.(\\ref{Pt}) with the inclusion of $\\eta\\neq 1$ in different emission models. Certainly the braking index varies when a pulsar evolves in this way. Further studies on this topic will be necessary and interesting.  There are a few pulsars which have initial periods $\\la 5$ ms in our calculation; for instance, 5 pulsars with $P_0<5$ ms appear in the right histogram of Fig.1 based on eq.(\\ref{a'}). This result can not be understood in the conventional opinion of $r-$mode instability,\\footnote{ The r-mode instability may spin down a nascent strange star or neutron star to a $P_0\\sim 3-5$ ms (e.g., Andersson \\& Kokkotas 2001). } % if it is not caused by statistical error or other factors being inherent in our $P_0$ calculations. However, as addressed in Xu \\& Busse (2001), the $r-$mode instability may be inhibited by differential rotation and strong magnetic field; the problem could thus be solved.  \\vspace{0.2cm} \\noindent {\\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). RXX wishes to sincerely acknowledge Mrs. S.Z. Peng for her help in preparing pulsar data two years ago."
    },
    "0206/astro-ph0206485_arXiv.txt": {
        "abstract": "The random error of radioastronomical measurements is usually computed in the weak-signal limit, which assumes that the system temperature is sensibly the same on and off source, or with and without a spectral line.  This assumption is often very poor.  We give examples of common situations in which it is important to distinguish the system noise in signal-bearing and signal-free regions. ",
        "introduction": "Few experiments are performed without some attempt at estimating their errors, and the random errors of measurement in radio astronomy are typically determined in one general way.  Some form of comparison is performed whereby samples are taken toward and away from a signal source, or with and without a spectral line. Subsequent analysis proceeds under the assumption that random errors everywhere in the dataset are as given by the statistical properties manifested in the signal-free regions.  No attempt is made to measure the variances of signal-bearing and signal-free samples separately during the experiment, and, after the fact, random errors of measurement in signal-bearing samples are obscured because the form of the signal is arbitrary. Discussions of fitting and profile analysis invariably assume that measurement variances are the same with or without the signal, as for instance the Zeeman analysis of \\cite{Mar95} or the fitting of functions ($e.g.$ Gaussians) by \\cite{KapSmi+66} or \\cite{Rie69}. Textbook discussions contain no suggestion that system noise is influenced by the presence of a signal or that samples with different variances may be interleaved in the same datastream \\citep{Kra86,BurGra01,RohWil00}.  Yet, such treatment has been flawed for a surprisingly long time. 100 K H I lines have been routinely observed with sub-100 K receiving systems for more than 30 years.  Continuum sources whose antenna temperatures exceed the equivalent noise temperature of the receiving equipment have been observed even longer.  The error of measurement in signal-bearing samples is often significantly different -- with current receivers it could easily be a factor of 5 at the peak of a strong galactic H I emission line -- but the difference has been ignored.  Error estimates determine confidence levels and even data containing strong signals can be compromised by misunderstanding of their significance; for instance, when two very strong signals are differenced to detect a smaller one in H I emission-absorption experiments and searches for Zeeman splitting. Considering how slowly experimental errors typically improve with the amount of time invested in an experiment, it follows that changes in the acknowledged errors of an experiment are equivalent to much larger differences in the observing time required to reach them. $A~priori$ knowledge of errors is an important element in the design of experiments and these considerations may have a significant effect on the planning of an observing session. They should be implemented in the software which supports analysis.  The purpose of this work is to illustrate a variety of common situations where random error is dominated by the presence of a signal.  In the following section some basics of radio astronomy measurement are sketched. These are used to analyse the statistics of noise and the errors of component fitting when signals are present in emission and absorption.  The final section is a brief summary with an even briefer mention of the extension of these notions to aperture synthesis. ",
        "conclusions": "Radio astronomers frequently observe signals which are strong enough to dominate the random errors of their experiments. Unfortunately, it is not always possible to recognize the effects which are induced and they are neglected.  Nonetheless, they have always been present in the data.  This discussion points up obvious deficiencies in extant data reduction software and analysis techniques.  Perhaps less obvious is the need not only for accurate calibration but also for reliable reporting on the part of the telescope systems.   Measurement errors cannot be accurately assessed and accomodated in downstream data handling unless the system, continuum and line antenna temperatures are preserved, along with knowledge of the mode of data-taking.  Synthesis instruments may be particularly difficult in this regard.  Consider the use of the VLA (say) to detect H I absorption against a continuum source at low galactic latitude in the presence of an emission profile like that shown in Fig. 1. The VLA does not return the total power or singledish spectra, or, equivalently, the variation of \\Tsys\\ across the passband.  The interferometer experiment {\\it per se} can only succeed to the extent that foreground emission disappears; only its added noise contribution remains.  We began the discussion by pointing out that the noise contributed from sky signals in single-dish observations occurs -- ignoring sidelobes, quantization noise and the like -- at those places and/or frequencies where the sources themselves are located.  It is an interesting endeavour to try to understand the extent to which source noise in interferometer experiments is similarly localized in the output datastream.  For phased arrays it would seem possible to reproduce the single-dish mode.  For synthesis arrays \\citep{AnaEke+89,CraNap89} the situation is much more complicated and uncertain even in the weak signal limit."
    },
    "0206/astro-ph0206216_arXiv.txt": {
        "abstract": "We report precise Doppler measurements of the stars HD~216437, HD~196050 and HD~160691 obtained with the Anglo-Australian Telescope using the UCLES spectrometer together with an iodine cell as part of the Anglo-Australian Planet Search.  Our measurements reveal periodic Keplerian velocity variations that we interpret as evidence for planets in orbit around these solar type stars. HD~216437 has a period of 1294$\\pm$250~d, a semi-amplitude of 38$\\pm$4 m~s$^{-1}$ and of an eccentricity of 0.33$\\pm$0.09. The minimum (M~sin~$i$) mass of the companion  is 2.1$\\pm$0.3~M$_{\\rm JUP}$ and the semi-major axis is 2.4$\\pm$0.5~au. HD~196050 has a period of 1288$\\pm$230~d, a semi-amplitude of 54$\\pm$8~m~s$^{-1}$ and an eccentricity of 0.28$\\pm$0.15. The minimum mass of the companion is 3.0$\\pm$0.5~M$_{\\rm JUP}$ and the semi-major axis is 2.3$\\pm$0.5~au. We also report further observations of the metal rich planet bearing star HD 160691. Our new solution confirms the previously reported planet and shows a trend indicating a second, longer-period companion. These discoveries add to the growing numbers of midly-eccentric, long-period extra-solar planets around Sun-like stars. As seems to be typical of stars with planets, both stars are metal-rich. ",
        "introduction": "Radial velocity programmes have now found around 80 extra-solar planets orbiting stars in the solar neighbourhood.  As the time baseline and precision of surveys improve new realms of possible planets are being explored. Discoveries include the first system of multiple planets orbiting a Sun-like star (Butler et al. 1999); the first planet seen in transit (Henry et al. 2000, Charbonneau et al. 2000); the first two sub-Saturn-mass planets (Marcy, Butler \\& Vogt 2000); and the Anglo-Australian Planet Searches' (AAPS) discovery of the first planet in a circular orbit outside the 0.1~au tidal-circularisation radius (Butler et al. 2001). The AAPS began operation in 1998, its southern hemisphere location completing all-sky coverage of the brightest stars at precisions reaching 3~m~s$^{-1}$. The AAPS has already found a number of extra-solar planets (Butler et al. 2001, 2002a; Tinney et al. 2001, 2002a; Jones et al. 2002). In this paper we present further results from this programme. ",
        "conclusions": "We report extra-solar planets in orbit around the stars HD~216437 and HD~196050, and further observations of HD~160691, which give the preliminary indication of a second planet. These detections serve to further emphasize that planetary systems with orbital parameters similar to those of our own Solar System are not as rare as suggested by the early extra-solar planet discoveries (e.g., Boss 2001). These discoveries confirm the preponderance (1) of relatively low-mass M~sin~$i$ planets and (2) planets around metal-rich objects. The detection of these relatively long-period planets gives us confidence in the stability of our search and gives added impetus for the continuation of the AAPS 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 (e.g. Marcy et al. 2002)."
    },
    "0206/astro-ph0206020_arXiv.txt": {
        "abstract": "We investigate the effects of thermonuclear reaction rate uncertainties on nova nucleosynthesis. One--zone nucleosynthesis calculations have been performed by adopting temperature--density--time profiles of the hottest hydrogen--burning zone (i.e., the region in which most of the nucleosynthesis takes place). We obtain our profiles from 7 different, recently published, hydrodynamic nova simulations covering peak temperatures in the range from T$_{peak}$=0.145--0.418 GK. For each of these profiles, we individually varied the rates of 175 reactions within their associated errors and analyzed the resulting abundance changes of 142 isotopes in the mass range below A=40. In total, we performed $\\approx$7350 nuclear reaction network calculations. We use the most recent thermonuclear reaction rate evaluations for the mass ranges A=1--20 and A=20--40. For the theoretical astrophysicist, our results indicate the extent to which nova nucleosynthesis calculations depend on presently uncertain nuclear physics input, while for the experimental nuclear physicist our results represent at least a qualitative guide for future measurements at stable and radioactive ion beam facilities. We find that present reaction rate estimates are reliable for predictions of Li, Be, C and N abundances in nova nucleosynthesis. However, rate uncertainties of several reactions have to be reduced significantly in order to predict more reliable O, F, Ne, Na, Mg, Al, Si, S, Cl and Ar abundances. Results are presented in tabular form for each adopted nova simulation. ",
        "introduction": "Classical novae occur in binary star systems consisting of a white dwarf and a main sequence star. When the companion star fills its Roche lobe, matter passes through the inner Lagrangian point and accumulates in an accretion disk before falling onto the white dwarf. The accreted layer gradually grows in mass. For sufficiently small mass--accretion rates, the deepest layers of the accreted material become partially degenerate. The temperature in the accumulated envelope increases because of compressional heating and energy release from nuclear reactions until a thermonuclear runaway occurs. At some time during the evolution, material from the white dwarf core is mixed into the accreted hydrogen--rich layer. As a consequence, a significant fraction of material, enriched in the products of hot hydrogen burning, is ejected into the interstellar medium. Spectroscopic studies of classical novae show enrichments of either C, N, O or of certain elements in the range from Ne to Ar (Gehrz et al. 1998, and references therein; Starrfield et al. 1998). The observed abundance patterns have been explained by assuming that the outbursts involve two fundamentally different types of white dwarfs with a composition consisting primarily of either carbon and oxygen (CO) or oxygen and neon (ONe).  The study of classical novae is of considerable interest for several reasons. First, spectroscopic studies of nova ejecta, when properly interpreted, reveal the composition of the underlying white dwarf, thereby constraining models of stellar evolution. Second, the observed elemental abundances also reflect the evolution of the thermonuclear runaway, such as peak temperatures and expansion time scales, and thus provide constraints for models of stellar explosions (Starrfield et al. 1998, 2000). Third, classical novae clearly contribute to the chemical evolution of the Galaxy. In fact, they have been proposed as the major source of the isotopes $^{13}$C and $^{17}$O, and perhaps $^{15}$N (Jos\\'{e} \\& Hernanz 1998). They may also represent a site for production of the cosmologically interesting isotope $^{7}$Li (Arnould \\& Norgaard 1975; Starrfield et al. 1978; Hernanz, Jos\\'{e}, Coc, \\& Isern 1996) as suggested by recent models of Galactic chemical evolution (Romano et al. 1999). Fourth, it is believed that radioactive isotopes are synthesized in nova outbursts. Short--lived isotopes, such as $^{14}$O ($\\tau$$_{1/2}$=71 s), $^{15}$O ($\\tau$$_{1/2}$=2 min) and $^{17}$F ($\\tau$$_{1/2}$=65 s), can reach the outer layers of the accreted envelope via convection, and their $\\beta$--decays provide an important energy source for the ejection of material (Starrfield et al. 1972). The decays of the short--lived nuclei $^{13}$N ($\\tau$$_{1/2}$=10 min) and $^{18}$F ($\\tau$$_{1/2}$=110 min) produce $\\gamma$--radiation of 511 keV and below, related to electron--positron annihilation and Compton--scattering, at a time when the expanding envelope becomes transparent to $\\gamma$--rays (G\\'{o}mez--Gomar et al. 1998; Hernanz et al. 1999). The decays of the longer--lived isotopes $^{7}$Be ($\\tau$$_{1/2}$=53 d) and $^{22}$Na ($\\tau$$_{1/2}$=2.6 y) produce $\\gamma$--rays with energies of E$_{\\gamma}$=478 and 1275 keV, respectively (Clayton \\& Hoyle 1974; Leising \\& Clayton 1987). Observations of $\\gamma$--rays from novae have been attempted with several satellites, but no positive detection has been reported. In the near future, however, novae will be promising targets for more sensitive instruments, such as the International Gamma--Ray Astrophysics Laboratory (INTEGRAL). Fifth, the discovery of $^{26}$Al ($\\tau$$_{1/2}$=7.4$\\times$10$^{5}$ y) in the interstellar medium (Mahoney et al. 1982) provided direct proof that nucleosynthesis is currently active in the Galaxy. >From the observed intensity of the 1809 keV $\\gamma$--ray line emission, it has been estimated that the production rate of $^{26}$Al in the Galaxy is $\\approx$2 M$_{\\odot}$ per 10$^{6}$ y. Although massive stars have been proposed as the main source of $^{26}$Al (Diehl et al. 1995; Prantzos \\& Diehl 1996; Diehl 1997; Kn\\\"{o}dlseder 1999), a contribution from classical novae cannot be ruled out (Politano et al. 1995; Jos\\'{e}, Hernanz, \\& Coc 1997). Sixth, the recent discovery of several presolar SiC grains with anomalous C, N, Al and Si isotopic ratios points towards a nova origin (Amari et al. 2001). If this identification is accurate, then the measured isotopic composition provides important constraints on both the nucleosynthesis and on the conditions in stellar outflows and circumstellar grain formation (Gehrz et al. 1998).  The thermonuclear runaway model reproduces several key features observed in nova outbursts. At present, the most successful calculations involve one--dimensional hydrodynamic codes that are directly coupled to large nuclear reaction networks (Kovetz \\& Prialnik 1997; Jos\\'{e} \\& Hernanz 1998; Starrfield et al. 1998, 2000; and references therein). However, some outstanding problems remain to be solved (Gehrz et al. 1998; Jos\\'{e} \\& Hernanz 1998; Starrfield et al. 1998, 2000). For example, the masses of the underlying white dwarfs are unknown and the rates of mass accretion are poorly constrained. The composition of white dwarfs involved in either CO or ONe novae is far from understood and may vary from outburst to outburst. The mechanism responsible for the mixing of white dwarf core material into the accreted hydrogen envelope is not universally accepted. The amount of mass ejected is controversial. Finally, many nuclear reaction cross sections entering in the hydrodynamic model calculations are uncertain by orders of magnitude.  In the present work, we focus on the effects of reaction rate uncertainties in nova model calculations. In the past, such effects were frequently ignored by stellar modelers who used only one specific set of recommended reaction rates from available libraries. Reaction rate uncertainties in hydrodynamic nova model calculations have been rarely explored in previous work. These studies were mainly concerned with the effects of a few uncertain reaction rates on the production of specific isotopes of interest, such as $^{18}$F (Coc et al. 2000), $^{22}$Na and $^{26}$Al (Jos\\'{e}, Coc, \\& Hernanz 1999), and Si--Ca (Jos\\'{e}, Coc, \\& Hernanz 2001). In the present work, we describe a more extensive approach. We independently vary the rates of 175 reactions that participate in nova model nucleosynthesis and analyze the resulting abundance variations of 142 isotopes in the mass range below A=40. In our calculations we take advantage of the two most recent thermonuclear reaction rate evaluations for the mass ranges A=1--20 (Angulo et al. 1999) and A=20--40 (Iliadis et al. 2001). For the theoretical astrophysicist, our results indicate the extent to which the nucleosynthesis depends on presently uncertain nuclear physics input, while for the experimental nuclear physicist our results represent at the least a qualitative guide for future measurements at stable or radioactive ion beam facilities.  Our philosophy and general issues related to the present work are described in $\\S$ 2. In $\\S$ 3 we explain our strategy and procedures in more detail. Results are presented in $\\S$ 4 and discussed in $\\S$ 5. A summary and conclusions are given in $\\S$ 6. ",
        "conclusions": "In the present work, we have investigated the effects of thermonuclear reaction rate uncertainties on nova nucleosynthesis. One--zone nucleosynthesis calculations have been performed by adopting temperature--density--time profiles of the hottest hydrogen--burning zone from 7 different, recent hydrodynamic nova simulations (Politano et al. 1995; Jos\\'{e} \\& Hernanz 1998; Jos\\'{e}, Coc, \\& Hernanz 1999; Starrfield et al. 2002). The adopted nova models cover peak temperatures in the range of T$_{peak}$=0.145--0.418 GK (Table 1). For each of these temperature--density--time profiles we have individually varied the rates of 175 reactions within their associated errors (Table 3) and analyzed the resulting abundance changes of 142 isotopes in the mass range below A=40. In total, we performed $\\approx$7350 reaction network calculations. We use the most recent thermonuclear reaction rate evaluations for the mass ranges A=1--20 (Angulo et al. 1999) and A=20--40 (Iliadis et al. 2001). Results are presented in tabular form for each adopted nova simulation (Tables 5--11). Figure 3 displays the results of reaction rate variations for a few selected cases. We find that present reaction rate estimates are reliable for predictions of Li, Be, C and N abundances in nova nucleosynthesis. However, uncertainties in the rates of several reactions have to be reduced significantly in order to predict more reliable O, F, Ne, Na, Mg, Al, Si, S, Cl and Ar abundances.  It is important to emphasize how to interpret the results of the present work. Hydrodynamic nova model calculations clearly show that typically only the outer layers of the envelope, not the deepest layers of the hydrogen--burning shell, are ejected after the thermonuclear runaway. The ejected layers are enriched, through convective mixing, with the products of the inner hydrogen--burning shell. From these considerations, it is clear that our calculations are unsuitable for defining {\\it absolute} isotopic abundances resulting from nova nucleosynthesis, since our one--zone calculations necessarily ignore convection ($\\S$ 2). Nevertheless, our procedure is adequate for exploring the effects of reaction rate uncertainties on abundance {\\it changes} in the hottest hydrogen--burning zone, i.e., the region in which most of the nucleosythesis takes place. It follows, therefore, that our final abundances (Table 4) should neither be compared to elemental abundances observed in nova ejecta nor to results from hydrodynamic model calculations. We also would like to stress the following point. If a particular reaction rate variation has insignificant effects on isotopic abundances in our calculations, then it is most likely that a full hydrodynamic model calculation will yield a similar result. However, the reverse statement is not neccessarily correct, i.e., if we find significant abundance changes as a result of a particular reaction rate variation, then a full hydrodynamic model calculation might not produce significant effects. Clearly, our work does not represent the final answer to the question of which reactions should be targets for future measurements, but should be regarded as a first step in that direction.  In Table 12 we summarize qualitatively some of our results. The table lists isotopes whose abundances change by more than a factor of 2 in at least one of the nova models considered here as a result of varying a particular reaction rate within uncertainties. It is striking that for the vast majority of reactions included in our network calculations, reaction rate variations have an insignificant effect on final isotopic abundances in all nova models. Instead, final abundances are influenced by variations of a restricted number of key reaction rates. Closer inspection of Tables 5--11 also shows that variations of the same reaction rates in nova models of the same white dwarf mass (e.g., models P1 and JCH2 with M$_{WD}$=1.25M$_{\\odot}$; or models P2 and S1 with M$_{WD}$=1.35M$_{\\odot}$) yield quantitatively different changes in final abundances. This is not surprising since different nova models assume different initial envelope compositions (Table 2) and achieve different peak temperatures (Table 1).  It can be seen from Table 12 and from Figure 3 that reaction rate variations of a few reactions, such as $^{23}$Na(p,$\\gamma$)$^{24}$Mg, $^{23}$Mg(p,$\\gamma$)$^{24}$Al, $^{30}$P(p,$\\gamma$)$^{31}$S and $^{33}$S(p,$\\gamma$)$^{34}$Cl, influence final abundances of a large number of isotopes. Consequently, new measurements of these reactions could significantly reduce uncertainties of isotopic abundances in nova model calculations. The reader might be surprised by the fact that certain reactions that were previously thought to play a role in nova nucleosynthesis do not appear in Table 12. In agreement with previous work (Iliadis et al. 1999), we find insignificant isotopic abundance changes as a result of $^{27}$Si(p,$\\gamma$)$^{28}$P, $^{31}$S(p,$\\gamma$)$^{33}$Cl, $^{35}$Ar(p,$\\gamma$)$^{36}$K and $^{39}$Ca(p,$\\gamma$)$^{40}$Sc reaction rate variations for all nova models. This result has been confirmed by recent hydrodynamic model calculations (Jos\\'{e} et al. 2001). The $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne and $^{19}$Ne(p,$\\gamma$)$^{20}$Na reactions, which were thought to cause a breakout of material from the CNO mass region to the region beyond Ne, are also missing in Table 12. Rate variations for both reactions have only small effects on final abundances in all nova models, except in model S1 which achieves the highest peak temperature (T$_{peak}$=0.418 GK). According to Table 7, an increase of those two reaction rates by a factor of 100 has only a moderate influence on abundance changes in the mass range below A=20. But even for this rather high peak temperature, no breakout of material from the CNO mass region is observed. This result has also been confirmed by recent hydrodynamic model calculations (Starrfield et al. 2002). It is also apparent from Tables 5--11 that ($\\alpha$,$\\gamma$) and ($\\alpha$,p) reactions in general are not important for nova nucleosynthesis.  Finally, it should be noted that it is difficult to estimate reliable reaction rate errors in certain cases. Consider as an example the $^{25}$Al(p,$\\gamma$)$^{26}$Si reaction. In this case, as for most other reactions involving short--lived target nuclei, we have assumed a reaction rate error of a factor of 100 up and down ($\\S\\S$ 3.3 and Table 3). An inspection of Tables 5--11 reveals only small abundance changes (within a factor of 2) as a result of varying the corresponding reaction rates within a factor of 100. However, for this particular case we have only limited experimental information regarding the energies of unobserved low--energy resonances (Iliadis et al. 1996). Depending on the location of these resonances, the $^{25}$Al(p,$\\gamma$)$^{26}$Si reaction rates could increase by much more than 2 orders of magnitude. As a consequence, the $^{26}$Al abundance will decrease significantly in all ONe nova models. Although not listed in Table 12, it is clear from this discussion that measurements of reactions such as $^{25}$Al(p,$\\gamma$)$^{26}$Si are also desirable in order to improve predictions of nova nucleosynthesis.    \\vspace{10mm}  The authors would like to thank A. Coc, M. Hernanz, R. Hix and M. Smith for stimulating discussions. We are also grateful for the detailed review of this work by the referee, S. Shore. This work was supported in part by the U. S. Department of Energy under Grant No. DE--FG02--97ER41041, by CICYT--PNIE ESP98--1348 and DGES PB98--1183--C03--02, and by Grants from NASA and NSF to ASU.   \\clearpage"
    },
    "0206/astro-ph0206284_arXiv.txt": {
        "abstract": "We report on two recent {\\em XMM-Newton} observations of Brown Dwarfs in the Pleiades cluster and in the field aiming to constrain the age dependence of X-ray emission from substellar objects. ",
        "introduction": "The standard picture of solar-type magnetic activity is expected to break down for very-low mass stars: being fully convective throughout the interior they lack the interface between radiative core and convective envelope in which the solar-type $\\alpha\\Omega$-dynamo is thought to reside. In spite of these theoretical predictions X-ray and H$\\alpha$ activity has been observed on stars with masses below the fully convective boundary, corresponding to spectral type $\\sim$\\,M3 (Fleming et~al. 1995, Gizis et~al. 2000). Recent H$\\alpha$ observations of ultracool field dwarfs indicate, however, a decline of activity setting in near the substellar limit at spectral type M9 (Mohanty \\& Basri 2002, see also Basri this volume).  While there is a substantial data base on chromospheric activity, only few X-ray observations of very low-mass (VLM) field dwarfs have been performed so far. Virtually all of the X-ray emitting field dwarfs with spectral type later than $\\sim$ M7 have been detected only during a temporary outburst, with quiescent emission below the detection threshold (see Sect.~3). Whether this is due to a lack of sensitivity of the respective observations, or whether these objects indeed are X-ray quiet can now be tested with a new generation of X-ray instruments onboard {\\em XMM-Newton} and {\\em Chandra}.  X-ray observations with {\\em ROSAT} have shown that young Brown Dwarfs (BDs) are more readily detected than older ones, as they show higher levels of activity (Neuh\\\"auser et~al. 1999). Due to the absence of an internal energy source the evolution of BDs goes along with a decrease of effective temperature. The accompanying drop of the ionization fraction may prevent coupling of the gas to the magnetic field, thus shutting off activity. Probing the relation between activity and the evolution of atmospheric conditions requires high-sensitivity X-ray observations of VLM stars and BDs at different ages. ",
        "conclusions": ""
    },
    "0206/astro-ph0206417_arXiv.txt": {
        "abstract": "We present high-resolution spectroscopy of the neutron-star/low-mass X-ray binaries 2S~0918$-$549 and 4U~1543$-$624 with the High Energy Transmission Grating Spectrometer onboard the {\\em Chandra X-ray Observatory} and the Reflection Grating Spectrometer onboard {\\em XMM-Newton}.  Previous low-resolution spectra of both sources showed a broad line-like feature at 0.7 keV that was originally attributed to unresolved line emission.  We recently showed that this feature could also be due to excess neutral Ne absorption, and this is confirmed by the new high-resolution {\\it Chandra\\/} spectra.  The {\\em Chandra} spectra are each well fit by an absorbed power-law $+$ blackbody model with a modified Ne/O number ratio of 0.52$\\pm$0.12 for 2S~0918$-$549 and 1.5$\\pm$0.3 for 4U~1543$-$624, compared to the interstellar-medium value of 0.18.  The {\\em XMM} spectrum of 2S~0918$-$549 is best fit by an absorbed power-law model with a Ne/O number ratio of 0.46$\\pm$0.03, consistent with the {\\em Chandra} result.  On the other hand, the {\\em XMM} spectrum of 4U~1543$-$624 is softer and less luminous than the {\\em Chandra} spectrum and has a best-fit Ne/O number ratio of 0.54$\\pm$0.03.  The difference between the measured abundances and the expected interstellar ratio, as well as the variation of the column densities of O and Ne in 4U~1543$-$624, supports the suggestion that there is absorption local to these binaries.  We propose that the variations in the O and Ne column densities of 4U~1543$-$624 are caused by changes in the ionization structure of the local absorbing material.  It is important to understand the effect of ionization on the measured absorption columns before the abundance of the local material can be determined.  This work supports our earlier suggestion that 2S~0918$-$549 and 4U~1543$-$624 are ultracompact binaries with Ne-rich companions. ",
        "introduction": "Low mass X-ray binaries (LMXBs) consist of a neutron star (NS) or black hole (BH) in orbit with a $\\lesssim$1~$M_{\\odot}$ companion.  An intriguing sub-class of LMXBs are the ultracompact binaries that have orbital periods less than 80 minutes.  Hydrogen-rich companions cannot sustain a LMXB system with such a short orbital period \\citep*{ps81,rjw82}.  However, orbital periods $\\lesssim$80 min are predicted for hydrogen-deficient or degenerate companions \\citep*{jar78,nrj86}, and this was confirmed with the orbital period measurements of the X-ray pulsar 4U~1626$-$67 \\citep[$P_{\\rm orb}$=42 min;][]{mmn+81}, the X-ray dipper 4U~1915$-$05 \\citep[$P_{\\rm orb}$=50 min;][]{ws82,wbm+82}, the X-ray bursters 4U~1820$-$30 \\citep*[$P_{\\rm orb}$=11 min;][]{spw87} and 4U~1850$-$087 \\citep[$P_{\\rm orb}$=21 min;][]{hcn+96}, and the detection of the white dwarf analogs, the AM CVn systems \\citep[e.g.,][]{w95}.  In addition, the three recently discovered millisecond X-ray pulsars XTE J1751$-$305, XTE J0929$-$314, and XTE J1807$-$294 were also found to be ultracompact binaries with $P_{\\rm orb}=$42, 44, and 35~min, respectively \\citep{mss+02,gcm+02,mss03}.  The conventional wisdom has been that the companions in ultracompact LMXBs are the remains of He white dwarfs (WDs) that have transferred a significant fraction of their mass to the NS.  However, recent X-ray spectral evidence indicates that some companions may be Ne-rich \\citep*{scm+01,jpc01}, which suggests the possibility of C-O or O-Ne-Mg WD companions.  The growing population of these exotic ultracompact systems, as well as the new evidence for Ne-rich companions, has interesting implications for the formation and evolution of binary systems.  Motivated by the observational evidence for Ne-rich donors, \\citet{ynh02} explored formation scenarios for these systems.  Here we discuss two other NS/LMXBs that may also be ultracompact binaries.  Both 2S~0918$-$549 ($l=275\\fdg9$, $b=-3\\fdg8$) and 4U~1543$-$624 ($l=321\\fdg8$, $b=-6\\fdg3$) have been observed by all of the major X-ray satellites since {\\it Uhuru}.  McClintock et al. (1978) identified the optical counterpart of 4U~1543$-$624 based on the {\\it SAS-3\\/} position, which was confirmed by {\\it HEAO 1}.  The flux measurements of 4U~1543$-$624 are roughly constant over the last 25 years \\citep{sak94,cs97,adn+00,jpc01,s02,ffm+03} with no periodicities from 50~s -- 10,000~s found in the {\\it EXOSAT\\/} data, and no periodicities from 0.1~s -- 1000~s detected in the {\\em SAX} data. Recently, \\citet{s02} presented spectral results from archival {\\em ASCA, SAX}, and {\\em RXTE} observations.  These results suggest that an increase in the luminosity of 4U~1543$-$624 is accompanied by a hardening of the spectrum.  In addition, an Fe-K emission line is seen in the hard state, but not detected in the low state.  \\citet{ffm+03} present an independent analysis of the two {\\em SAX} observations which show no significant luminosity change but a spectral hardening in the second observation.  Both \\citet{ffm+03} and \\citet{adn+00} report Fe-K line detections for the {\\em SAX} and {\\em ASCA} observations, respectively.  In contrast, 2S~0918$-$549 shows a factor of 10 X-ray variability, but again has no known periodicities in either the X-ray or optical bands down to timescales of 1 hour (Forman et al. 1978; Warwick et al. 1981; Chevalier \\& Ilovaisky 1987; Smale \\& Lochner 1992; Christian \\& Swank 1997; Schulz 1999; Jonker et al. 2001).  Chevalier \\& Ilovaisky (1987), using the {\\it Einstein\\/} HRI position, identified an ultraviolet-bright optical counterpart for 2S~0918$-$549 and suggested a source distance of 15 kpc based on the properties of other LMXBs, i.e., $M_{V}$$=$0.0 and $(B-V)_{0}$$=$0.0.  Recently, Jonker et al. (2001) detected a type I X-ray burst from 2S~0918$-$549 and derived an upper limit to the distance of 4.9 kpc~from radius expansion arguments.  A similar analysis using {\\em BeppoSAX} Wide Field Camera data of 2S~0918$-$549 implied a distance of 4.2~kpc \\citep{cvz+02}.  We have identified both 2S~0918$-$549 and 4U~1543$-$624 as being part of a class of four NS/LMXBs all having a similar feature at 0.7~keV in their low-resolution spectra \\citep{jpc01}.  This feature had been attributed to unresolved line emission from Fe and O \\citep*[see, e.g.,][]{cws94,wka97}.  However, a high-resolution observation of the brightest of these sources, 4U~0614$+$091, with the {\\em Chandra X-Ray Observatory}, failed to detect any emission lines, finding instead an unusually high Ne/O number ratio in the absorption along the line of sight (Paerels et al. 2001).  Previously, we showed that the {\\em ASCA} spectra of all four sources are well fit {\\em without} a 0.7~keV emission line using a model that includes photoelectric absorption due to excess Ne along the lines of sight and presumably local to the sources.  In this paper, we present the high-resolution {\\em Chandra} and {\\em XMM-Newton} spectra of two sources, 2S~0918$-$549 and 4U~1543$-$624, which support our earlier {\\em ASCA} results.  Given the high $L_{\\rm X}/L_{\\rm opt}$ ratio of these sources and the excess Ne absorption, we suggested that the systems are ultracompact binaries with Ne-rich degenerate donors \\citep{jpc01}.  This is similar to the ultracompact X-ray pulsar 4U~1626$-$67.  We note that while the spectrum of 4U~1626$-$67 does show O and Ne emission lines \\citep{awn+95,scm+01}, it does not resemble the sources we are considering in either low or high-resolution spectra. While we do compare our results to those of 4U~1626$-$67, it is important to consider the differences between the systems when drawing conclusions. ",
        "conclusions": "We have shown that the {\\it Chandra}/HETGS and {\\em XMM}/RGS spectra of 2S~0918$-$549 and 4U~1543$-$624 are well fit by models that allow for absorption columns of neutral Ne and O with abundance ratios significantly different from the expected ISM ratio.  The {\\em Chandra} spectrum of 2S~0918$-$549 has a best fit Ne/O number ratio of 0.52$\\pm$0.12, or roughly 3$\\times$ the ISM value of 0.18.  The {\\em XMM} fit of 2S~0918$-$549 gives a consistent Ne/O number ratio of 0.46$\\pm$0.03.  The best-fit model for the {\\em Chandra} spectrum of 4U~1543$-$624 has a Ne/O number ratio of 1.5$\\pm$0.3, while the best-fit {\\em XMM} model has a Ne/O number ratio of 0.54$\\pm$0.03. The unusual abundance ratios as well as the variability seen in the observations of 4U~1543$-$624 lead us to conclude that there is absorption local to these binaries and that the material is Ne enriched.  From the $L_{\\rm X}/L_{\\rm opt}$ ratios of 2S~0918$-$549 and 4U~1543$-$624, we expect these binaries to have orbital periods $\\lesssim$$60$~min based on the empirical relationship determined by van Paradijs \\& McClintock (1994).  We searched the {\\it Chandra\\/} lightcurves for orbital modulation, but found no modulations larger than 1--2\\%.  {\\em XMM} provided a more sensitive search for orbital modulations, but we found no signatures of orbital modulation at fractional rms upper limits of 0.14--0.3\\%.  Confirmation of such short orbital periods would place 2S~0918$-$549 and 4U~1543$-$624 in the class of ultracompact LMXBs (P$_{\\rm orb}$$\\lesssim$80 min). Ultracompact binaries require H-depleted or degenerate dwarf companions (Joss et al. 1978; Nelson et al. 1986).  Such companions would be expected to have non-standard abundances compared to ISM values.  The {\\it Chandra\\/} spectrum of 4U~1626$-$67 revealed absorption edges of C, O, and Ne which are 5 times larger than would be predicted given the hydrogen column density measured in the UV. For 4U~1626$-$67, the measured local abundance ratios, if the excess material is assumed to have originated around the binary, are consistent with the expected abundances in the chemically fractionated core of a C-O or O-Ne-Mg WD (Schulz et al. 2001).  The evidence for Ne absorption in 2S~0918$-$549 and 4U~1543$-$624 hints at a similarity between these sources and 4U~1626$-$67.  In addition, there is evidence that the optical spectra of this group of sources is also similar to 4U~1626$-$67 with no H or He lines detected, but with a \\ion{C}{3}/\\ion{N}{3} emission line near 4640~\\AA\\/ (Wang \\& Chakrabarty 2003, in preparation).  Based on the similarity between these sources and 4U~1626$-$67, we previously attributed the excess Ne absorption in 2S~0918$-$549 and 4U~1543$-$624 to material local to the sources, and suggested that these systems contained a Ne-rich degenerate donor \\citep{jpc01}.  The one assumption we have made in this analysis is that the absorption is from neutral material only.  If there is a sizable contribution to the absorption from material local to the binary, this assumption is most likely not valid, since we would expect ionization of the local material from the central source.  Accounting for ionization of the material may explain the large Ne/O ratios observed. We find Ne/O ratios larger than the 0.22 inferred for the local absorption in the {\\em Chandra} spectrum of 4U~1626$-$67 \\citep{scm+01}.  Ne/O ratios for local absorption in our source will be even higher than the measured total line-of-sight ratio once ISM contributions are removed, since the O column from the ISM should be larger than the Ne column. Unfortunately, we can not do this calculation, as was possible for 4U~1626$-$67, since we do not know the expected ISM contribution for 2S~0918$-$549 or 4U~1543$-$624.  If the local material is affected by ionization, the effect will be different for each element.  O will become ionized before Ne, leading to an enhanced Ne/O ratio as measured by the neutral edges.  While it is tempting to assume that ionization is the only cause for the unusual abundances, we point out that if the local material was of standard abundances, we would expect local absorption from higher Z elements, like Mg and Si, which are not seen in the spectra of 2S~0918$-$549 or 4U~1543$-$624.  This leads us to conclude that material must have some enhancement of Ne to show such strong absorption.  One of the most interesting results is the difference in the Ne and O column densities in the two observations of 4U~1543$-$624.  The {\\em XMM} results show a increase in the O column density with a decrease in the Ne column density.  In addition, the {\\em XMM} spectrum is both softer and has a lower luminosity in the 2--10~keV band than the {\\em Chandra} spectrum, which suggests that the source was not in the same state in the {\\em Chandra} and {\\em XMM} observations.  We suggest that the difference in the 4U~1543$-$624 results provides support for the local nature of the absorption.  Since the high-resolution results of 2S~0918$-$549 are consistent not only in the measured Ne/O number ratio but also in the spectral model in general, we are confident that instrumental differences do not have a significant effect on the spectral results.  It is interesting to note the the Ne/O ratios found in the {\\em Chandra} and {\\em XMM} observations of 2S~0918$-$549 are smaller than the inferred value from the {\\em ASCA} result of \\citet{jpc01}, while at the same time the flux of 2S~0918$-$549 has decreased by a factor of 10 since the {\\em ASCA} observation.  It is possible that the state changes alter the ionization structure of the local material, which in turn changes the absorbing columns of neutral Ne and O.  It will be difficult to separate the effect of unusual abundances from ionization changes without more data on these sources. We propose that multiple observations of 4U~1543$-$624 using simultaneous high-resolution and broadband observations will probe the connection between spectral state differences and the column density variations.  It is important to note that ionization effects and the ISM contribution must be understood before it is possible to determine the intrinsic Ne/O abundance ratio of the local material, which can then be used to place constraints on the composition of the companion.  Although we attribute the excess Ne absorption in 2S~0918$-$549 and 4U~1543$-$624 to local material in both of the binaries, we also consider the possibility that the excess Ne is due to enhancements of the ISM along the line of sight.  The best measurement of absorption toward an X-ray binary, the {\\it Chandra\\/} spectral analysis of Cyg X-1, shows columns of O and Ne that are consistent with standard ISM abundances (Schulz et al. 2002).  To quantify the variations of the ISM abundances, we have undertaken a study of the ISM using column density measurements from the spectra of X-ray binaries (Juett et al. 2003).  Initial results show that other LMXBs do not show the Ne/O ratios measured for these sources.  At most, Ne/O number ratios are only twice the expected ISM values.  This is still significantly smaller than we find in 2S~0918$-$549 and 4U~1543$-$624, but does suggest that there are some variations between lines-of-sight.  We still favor the interpretation that there is absorbing material local to the binaries, and that this is strengthened by the observation-to-observation variations in the Ne/O number ratio of 4U~1543$-$624 which could not be due to the ISM.  Analysis of the {\\it Chandra\\/} and {\\em XMM} spectra of 2S~0918$-$549 and 4U~1543$-$624 found no Ne or O lines like those seen in 4U~1626$-$67.  The absence of strong lines in the spectra of 2S~0918$-$549 and 4U~1543$-$624 demonstrates that their circumbinary environments are different compared to 4U~1626$-$67.  One possible reason for this difference may be that 4U~1626$-$67 is a pulsar with $B$$=$$3\\times10^{12}$~G (Orlandini et al. 1998), while 2S~0918$-$549 and 4U~1543$-$624 are likely to have weak magnetic fields ($\\sim$10$^8$ G).  This undoubtedly results in a substantial difference in the accretion flow geometry.  There is weak evidence for a line-like feature at 6.4~keV in the {\\em XMM} spectrum of 4U~1543$-$624.  This feature is very broad and could possibly be due to a mis-modeling of the continuum emission rather than a Fe line.  In addition, there is evidence for structure around the Ne edge in the spectra of both sources.  It is likely that this is due to instrumental effects but could also be due to structure in either neutral Ne edges or possibly due to confusion with ionized O edges.  Recent {\\it RXTE\\/} and {\\em BeppoSAX} observations of 2S~0918$-$549 found the first thermonuclear X-ray bursts from this source \\citep{jvh+01,cvz+02}.  Both bursts were short, with durations $\\lesssim$$100$~s.  These short bursts suggest that H and/or He is undergoing unstable burning on the surface of the NS.  If the companion has a large C abundance, we might expect to see a much longer burst ($\\approx$1 hr), like the ``superbursts'' seen in some LMXBs, which have been attributed to thermonuclear burning of C on neutron star surfaces (Cornelisse et al. 2000; Cumming \\& Bildsten 2001; Strohmayer \\& Brown 2001).  The more ordinary properties of the X-ray bursts from 2S~0918$-$549 may indicate that the donor is not a C-O dwarf as suggested for 4U~1626$-$67.  In this case, another mechanism for Ne enhancement is needed.  On the other hand, it may still be possible for a carbon accretor to show short X-ray bursts. One possibility is that while the companion is H-deficient, there is still a non-negligible H fraction that is accreted by the NS and is then responsible for the Type I X-ray bursts \\citep[e.g.,][]{nrj86}. Alternatively, the heavy elements (C, O, and Ne) may undergo spallation during accretion, leaving He and H nuclei which could then undergo unstable thermonuclear burning in the usual way (Bildsten, Salpeter, \\& Wasserman 1992).  Spallation reactions would produce $\\gamma$-ray emission lines at 4.4 and 6.1 MeV.  Unfortunately, the strength of these lines as calculated by Bildsten et al. (1992) is below current observational detection limits.  For a 1~Ms observation with {\\em Integral}, the detection sensitivities are $\\sim$1000 times higher than the most optimistic line flux estimates."
    },
    "0206/nucl-ex0206002_arXiv.txt": {
        "abstract": "The $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne reaction is one of two routes for breakout from the hot CNO cycles into the $rp$ process in accreting neutron stars. Its astrophysical rate depends critically on the decay properties of excited states in $^{19}$Ne lying just above the $^{15}$O + $\\alpha$ threshold. We have measured the $\\alpha$-decay branching ratios for these states using the $p(^{21}$Ne,$t)^{19}$Ne reaction at 43 MeV/u. Combining our measurements with previous determinations of the radiative widths of these states, we conclude that no significant breakout from the hot CNO cycle into the $rp$ process in novae is possible via $^{15}$O($\\alpha$,$\\gamma$)$^{19}$Ne, assuming current models accurately represent their temperature and density conditions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206147_arXiv.txt": {
        "abstract": "{We combine near-infrared photometry obtained with the VLT/ISAAC instrument and archival HST/WFPC2 optical images to determine $VIK$ magnitudes and colours of globular clusters in two early-type galaxies, NGC~3115 and NGC~4365. The combination of near-IR and optical photometry provides a way to lift the age-metallicity degeneracy. For NGC 3115, the globular cluster colours reveal two major sub-populations, consistent with previous studies. By comparing the $V-I$, $V-K$ colours of the NGC~3115 globular clusters with stellar populations models, we find that the colour difference between the two $\\ga10$ Gyr old major sub-populations is primarily due to a difference in metallicity. We find $\\Delta$[Fe/H]$=1.0\\pm0.3$ dex and the blue (metal-poor) and red (metal-rich) globular cluster sub-populations being coeval within 3 Gyr. In contrast to the NGC~3115 globular clusters, the globular cluster system in NGC~4365 exhibits a more complex age and metallicity structure. {\\it We find a significant population of intermediate-age very metal-rich globular clusters} along with an old population of both metal-rich and metal-poor clusters. Specifically, we observe a large population of globular clusters with red $V-K$ colours but intermediate $V-I$ colours, for which all current stellar population models give ages and metallicities in the range $\\sim2-8$ Gyr and $\\sim0.5Z_\\odot- 3Z_\\odot$, respectively. After 10 Gyr of passive evolution, the intermediate-age globular clusters in NGC~4365 will have colours which are consistent with the very metal-rich population of globular clusters in giant elliptical galaxies, such as M87. Our results for both globular cluster systems are consistent with previous age and metallicity studies of the diffuse galactic light. In addition to the major globular cluster populations in NGC~3115 and NGC~4365 we report on the detection of objects with extremely red colours ($V-K\\ga3.8$ mag), whose nature could not ultimately be revealed with the present data.} ",
        "introduction": "\\label{ln:intro} Globular cluster systems are useful tracers of galaxy evolution. They consist of bright clusters, which are made of stars that share the same age and chemical composition. Globular clusters can form during major merger events in galaxies \\citep[e.g.][etc.]{holtzman92, whitmore93, whitmore95, schweizer96}, but we also observe the formation of massive star cluster in galaxies with a moderate star formation rate \\citep[e.g.][etc.]{oconnell94, barth95, oconnell95, brandl96, larsen99, hunter00}. In recent years the study of globular cluster systems revealed the presence of globular cluster sub-populations which must have formed in multiple formation epochs and/or mechanisms \\citep[see][for reviews]{ashman98, kisslerpatig00, vandenbergh00, harris01}.  In practice, for distant galaxies, we observe integrated properties of their globular clusters, and would like to deduce their physical properties, in particular their age and metallicity. Unfortunately, even in systems where reddening can be taken as uniform, the optically observed parameters suffer from the well known age-metallicity degeneracy. Spectroscopy can overcome this infamous problem \\citep{jones95, worthey97, vazdekis99a} to a large extent, but is, however, very time consuming to perform for hundreds of objects. Photometry still represents the most efficient way to study an entire globular cluster system.  There have been several attempts to solve this degeneracy for globular cluster systems in early-type galaxies which employed optical colours only \\citep[e.g.][]{kisslerpatig97, whitmore97, kisslerpatig98, kundu99, puzia99} with various degrees of success. Previous studies used the optical {\\it colour difference and the turn-over magnitude difference} between the major globular cluster sub-populations to derive an age and metallicity difference. This method suffers from two main problems: 1) It is based on the assumption that the mass functions are similar for both sub-populations, so that luminosity differences directly reflect differences in the mass-to-light ratios of the individual populations. 2) The peak colour and turn-over magnitude differences are small and somewhat model dependent.  Alternatively one could derive mean age and metallicity differences between globular cluster sub-populations from {\\it colour-colour diagrams}, which does not depend on the globular-cluster mass functions.  The combination of optical and near-infrared (near-IR) photometry can largely reduce the age-metallicity degeneracy. This is because optical to near-IR indices like $V-K$ are very sensitive to metallicity, but only have modest age sensitivity \\citep[hereafter Paper I]{kisslerpatig00, puzia01, kisslerpatig02}. Physically, this technique works because the $V$-band samples mainly the light of stars near the turn-off, while the $K$-band is most sensitive to cooler stars on the giant-branch in old stellar populations \\citep[see e.g.][]{yi01}. While the turn-off is mostly affected by age, the giant branch is primarily sensitive to metallicity \\cite[e.g.][]{saviane00}. Thus, $V-I$ and $V-K$ are affected similarly by age, but $V-K$ is much more sensitive to metallicity (see Table \\ref{tab:slopes} below). This allows the age-metallicity degeneracy to be lifted by determining the location of a simple stellar population in a plot of $V-I$ vs. $V-K$.  Such diagrams are powerful tools to study the age and the metallicity of globular-cluster populations. Including near-IR passbands to our optical data we can increase the sensitivity to metallicity differences by a factor of two or more (see Paper I, and SSP models of e.g. \\citealt{bc00}).  \\begin{table}[t!] \\centering \\caption[width=\\textwidth]{Basic information on observed galaxies. The references are (1): \\cite{RC3}, (2): \\cite{schlegel98}, (3): \\cite{buta95}, (4): \\cite{frogel78}. (5): \\cite{tonry01}. } \\label{tab:galdat} \\begin{tabular}{l r r l} \\hline \\noalign{\\smallskip} Parameter & NGC~3115 & NGC~4365 & Ref. \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} type               & S0             & E3              &(1) \\\\ RA  (J2000)        &10h 05m 14s     & 12h 24m 28s     &(1) \\\\ DEC (J2000) &$-07^{\\rm o}$ 43' 07'' & $+07^{\\rm o}$ 19' 03'' &(1)\\\\ $l$                &$247.78^{\\rm o}$&$283.80^{\\rm o}$ &(1)\\\\ $b$                &$36.78^{\\rm o}$ &$ 69.18^{\\rm o}$ &(1)\\\\ $B_{\\rm T,0}$      & 9.74           & 10.49           &(1) \\\\ E$_{B-V}$          & 0.047          &  0.021          &(2)\\\\ $(B-V)_{\\rm o}$    & 0.94           &  0.95           &(1)\\\\ $(V-I)_{\\rm eff,o}$& 1.25           &  1.25           &(3)\\\\ $(V-K)_{\\rm eff,o}$&$3.30\\pm0.02$   & $3.29\\pm0.1$    &(4)\\\\ $(m-M)_V$          &$29.93\\pm0.09$  & $31.55\\pm0.17$  &(5)\\\\ $M_V$              &$-21.13$        &$-22.01$         &(1),(5)\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table}  In this paper we study the globular cluster systems of NGC~3115 and NGC~4365. NGC 3115 is an isolated galaxy, located at the very tip of the southern extension of the Leo Group, with just one significant accompanying nucleated dwarf elliptical galaxy \\citep{puzia00a}. It features a bimodal optical colour distribution of globular clusters \\citep{kundu98,gebhardt99,larsen01}.  NGC~4365 is a cluster giant elliptical galaxy located at the outer edge of the Virgo Cluster. Optical photometry of the globular cluster system revealed a broad but single-peak colour distribution with little evidence for bimodality \\citep{forbes96, gebhardt99, larsen01, kundu01a}. \\citeauthor{surma95} (\\citeyear{surma95}, see also \\citealt{davies01}) detected a decoupled, counter-rotating core, which consists of a younger and more metal-rich stellar population than the rest of the galaxy. Such decoupled cores are believed to be created in major merger events. All relevant galaxy properties are summarized in Table \\ref{tab:galdat}.  The major goal of this paper is to derive ages and metallicities of globular clusters from a comparison of optical/near-IR colours with several SSP models. Putting this results into a larger context, including previous findings from galactic integrated light studies, we can further constrain formation scenarios for both galaxies.  The outline of this paper reads as follows. In Section 2 we describe the data reduction and calibration of our new near-IR VLT/ISAAC and the archival optical HST/WFPC2 data. The main results are presented in Section 3 followed by a discussion in Section 4. All major findings are summarized in Section 5. ",
        "conclusions": "\\label{ln:discussion} \\subsection{NGC~4365, a Cooling-Down Galaxy} \\cite{davies01} find a luminosity-weighted age of $\\sim12$ Gyr for the solar-metallicity kinematically-decoupled core (KDC) in NGC~4365 using the diagnostic plot (H$\\beta$ vs.~[MgFe]) and SSP models of \\cite{vazdekis99}. From the comparison of our near-IR/optical colours with the same models, we find that $\\sim40-80$\\% of our sample are metal-rich ($Z\\ga0.5\\, Z_\\odot$) globular clusters with ages between $\\sim2-8$ Gyr. However, within our observational measurement errors the age distribution, derived from the linear interpolation of SSP models, is consistent with a single-burst of globular-cluster formation which took place $\\sim3-4$ Gyr ago. Together with the metal-rich globular clusters a part of the stellar population in the KDC must have been formed from the same pre-enriched gas in a significant star-formation event. In fact, \\cite{davies01} find that the $\\sim2$ Gyr younger luminosity-weighted age of the KDC (compared to the outer parts of the galaxy) can be accounted for by a contaminating, 5 Gyr old stellar population of solar metallicity. This young population would amount 6\\% of the mass inside $r\\la1.6$\\arcsec\\ or $\\la160$ pc.  It is tempting to associate the intermediate-age globular cluster population with the kinematically decoupled core. Within the prediction accuracy of SSP models this young KDC population is consistent with being the diffuse counterpart of the stellar population in metal-rich intermediate-age globular clusters. Numerical simulations of gaseous mergers \\cite[see e.g.][for a review]{barnes92} predict the settling of the progenitors' gas into the inner parts of the merger remnant within $\\sim10^8$ years. If globular clusters and the central stellar population were in fact formed from the same gas, both stellar populations have to be almost coeval.  We can compare the stellar mass contained in an intermediate-age component in NGC~4365 with the mass of the intermediate-age globular cluster population to test how consistent these are with forming in the same event. To calculate the mass of the diffuse stellar component, we begin with the decoupled core. \\citeauthor{davies01} found that the intermediate-age component makes up 6\\% of the stellar mass of the decoupled core. Given the total stellar mass estimate for the core of $M_{\\rm core}\\la (8.1\\pm2.5)\\cdot10^9 M_\\odot$ from \\cite{surma95}, the resulting mass in the core of the intermediate-age diffuse stellar component is $(4.9\\pm2.5)\\cdot10^8 M_\\odot$.  This mass is a lower limit to the total mass of the diffuse intermediate-age stellar population because it is unlikely that this component is found only in the very central region. In particular, the intermediate-age globular clusters are observed over a fairly extended region, and it would be difficult to arrange for the corresponding diffuse component not to have a similar extended spatial distribution. Therefore, in order to estimate the total stellar mass of the galaxy at intermediate ages, we need to estimate the fraction of the total galaxy mass coming from the intermediate age population. This fraction can not be greater than a few percent, or the extended intermediate-age population would have been detected by \\citeauthor{davies01} as they detected such a component in the central region. Conversely, the mass fraction of the intermediate-age component must be at least a few-tenths of a percent based on comparing the intermediate-age stellar mass of $(4.9\\pm2.5)\\cdot10^8 M_\\odot$ observed in the core to the total galaxy stellar mass of $\\sim4.8\\cdot10^{11}M_\\odot$ \\citep{bender89}.  To bracket these possibilities, we carry out a representative calculation with a 1\\% fraction of the overall stellar mass of the galaxy in an intermediate-age component, and show how the results are affected by varying this in either direction. This fraction gives a stellar mass over the whole galaxy in the intermediate-age component of $\\sim5.2\\cdot10^9M_\\odot$.  What is the corresponding mass currently in the intermediate-age GC system? After correcting our sample for incompleteness in luminosity and spatial coverage and with the assumption that $\\sim40-80$\\% of our sample globular clusters have intermediate-ages, we find that the total number of intermediate-age globular clusters in NGC~4365 is $\\sim 150$ (within a factor of 2). To determine the mass in this system, we adopt the two-component mass function of \\cite{mclaughlin96} over the mass range of $10^4-10^8 M_\\odot$, appropriate for the GCs of giant elliptical galaxies. This gives a total mass of $\\sim1.6\\cdot10^8M_\\odot$ in the globular cluster system. A similar result is found by simply counting up the estimated mass of the observed clusters and accounting for the spatial areas not covered by our observations.  The efficiency of formation and survival to several Gyr of the intermediate-age globular cluster system is 3\\%. This follows directly from the derived mass of $\\sim1.6\\cdot10^8M_\\odot$ for the intermediate-age globular cluster system and the corresponding $\\sim5.2\\cdot10^9M_\\odot$ for the diffuse stellar component of the intermediate-age system. If the adopted mass fraction of the galaxy is too high, the formation and survival efficiency increases to rather high values of 10\\% or more. If the adopted mass fraction of the galaxy in intermediate-age component is a factor of a few too low, this number drops to about 1\\%. It is difficult to make the data consistent with a formation efficiency less than 1\\%.  How do these formation and survival efficiencies compare to other observations? Observations of strongly starbursting systems show that $\\sim$ 20\\% of the stars form in compact star clusters \\citep[e.g.][]{meurer95, whitmore99, larsen99, zepf99}. This represents the peak formation efficiency of globular clusters. Studies of older mergers find that a few percent of the stars are in globular clusters \\citep[e.g.][]{schweizer96, miller97, zepf99, goudfrooij01b}. The best determined case for older globular cluster systems is probably the Galactic halo, in which about 1\\% of the stars are in globular clusters. A global average of all spheroidal galaxy systems is around 0.3\\% \\citep[e.g.][]{mclaughlin99}. However, there are large variations between galaxies, and within sub-populations in individual galaxies. Dynamical evolution may significantly reduce the number of globular clusters observed at old ages compared to initial values at young ages \\citep[e.g.][]{gnedin97, vesperini00}.  The comparison of the formation and survival efficiency of the NGC~4365 intermediate-age globular cluster system to other systems suggests that NGC~4365 and older mergers have about 2\\% of their stars in globular clusters, with an uncertainty of a factor of a few. It is possible that the stellar mass attributed to the younger component should be restricted to that in the core alone, in which case the formation and survival efficiency of the NGC~4365 globular cluster system would be more like the peak efficiency seen in recent mergers with little subsequent dynamical evolution. Underestimating the mass in the diffuse intermediate-age population would reduce the formation and survival efficiency by a factor of a few, but it is difficult to make the data consistent with globular cluster formation efficiencies less than a percent or so.  We can also compare the metallicity of the intermediate-age globular cluster and diffuse stellar components. NGC~4365 shows super-metal-rich globular clusters ($Z>2\\,Z_\\odot$) while the luminosity-weighted stellar light of the decoupled core reveals a roughly solar metallicity \\citep{surma95, davies01}. This is likely due to a mixing effect of the integrated light of stellar populations in the core. If stars and globular clusters were formed from the same pre-enriched material, super-solar metallicities are expected in the core, as well. The small fraction of light from super-metal-rich stars is likely to be washed out by insufficient resolution of the inner structures. With a spatial resolution of $\\sim1.4$\\arcsec\\ (radial size of the decoupled core $\\sim8$\\arcsec ) \\cite{surma95} report that the metallicity is at least $2.3Z_\\odot$ in the center of NGC~4365. At face value, the core metallicity and the metallicity of the most metal-rich globular clusters are consistent.  \\subsection{NGC~3115, a Quiescent Galaxy} It seems that the colour difference between the two major sub-populations of globular clusters in NGC~3115 is mainly due to a difference in metallicity, rather than in age. Within the errors this result is still consistent with the findings of \\cite{kundu98} who report, using optical photometry only, that the metal-rich clusters are about $\\sim3$ Gyr younger than the metal-poor ones.  In contrast to NGC~4365, NGC~3115 seems to have experienced no major star-formation event in the last $\\sim$10 Gyr. Both its globular cluster sub-populations are roughly that old or older. This resembles the situation in the globular cluster system of the Milky Way, where the two sub-systems, corresponding to thick-disk/bulge and halo populations, differ mainly in metallicity \\citep[][etc.]{minniti95, cote00} with a relatively small age difference where the metal-rich clusters may be slightly younger than their metal-poor counterparts \\citep{rosenberg99}.  Quantitatively, in NGC~3115 it remains to be seen whether or not the specific frequency $S_{\\rm N}$ (number of globular clusters per unit luminosity, see \\citealt{harris81}) for the metal-poor stellar populations (i.e., the ratio of metal-poor globular clusters to metal-poor halo stars) is also roughly three times higher than for the metal-rich stellar populations, as it is observed in NGC~5128 and M31.  Given the limited information about the halo stars, we conclude that the globular cluster system and the majority of halo stars in NGC~3115 most likely formed in two epochs and/or by two different mechanisms more than $\\sim$10 Gyr ago. Ever since then, the stellar populations in globular clusters and in the galaxy halo seem to have undergone a purely passive evolution."
    },
    "0206/astro-ph0206371_arXiv.txt": {
        "abstract": "Upper limits are presented on the diffuse flux of ultra-high energy $\\nu_e$, based on analysis of data taken by the RICE experiment during August, 2000. The RICE receiver array at South Pole monitors cold ice for radio-wavelength Cherenkov radiation resulting from neutrino-induced in-ice showers. For energies above 1 EeV, RICE is an effective detector of over 15 km$^3$ sr. Potential signal events are separated from backgrounds using vertex location, event reconstruction, and signal shape. These are the first terrestrial limits exploiting the physics of radio Cherenkov emissions from charged-current $\\nu_e+N\\to e+N'$ interactions. ",
        "introduction": "\\label{intro.sec}  Neutrinos are generally expected to be an important component of ultra-high energy (UHE) cosmic rays\\cite{GaisserHS}. Neutrinos propagate in straight lines from their source and retain spectral information about their origin - in contrast to protons, electrons or gamma-rays, which interact with cosmic magnetic fields, cosmic radiation backgrounds at microwave and radio frequencies, and material within the source itself. Energies below 100 GeV are the province of solar, supernova and atmospheric neutrino experiments. Above 100 GeV it is expected that astrophysical sources (e.g. AGN, GRBs, the galactic disk) will be observable in planned cubic kilometer scale detectors\\cite{IceCube,ANTARES,NEMO,NESTOR}, or perhaps in prototype existing experiments that have pioneered the development of large arrays of optical Cherenkov detectors\\cite{AMANDA,BAIKAL}. Despite their impressive size, even these experiments may be too small to see scientifically interesting fluxes at energies above 10 PeV. Of particular interest, neutrinos of energy 0.1-1 EeV may provide key information for understanding anomalies reported in the cosmic ray spectrum\\cite{AGASA,HPARK,FLYSEYE} above the ``GZK\" cutoff\\cite{GZK}. Detection of the nominal flux of GZK neutrinos\\cite{BerZat69} requires a detector mass equivalent to 100 km$^3$ of water\\cite{EngelSS}. To achieve such large masses, other techniques have been proposed for this energy range, including detection of horizontal air showers\\cite{AUGER,CapelleCPZ,OWL,EUSO}, acoustic signals in water or ice\\cite{acoustic} and a variety of ideas\\cite{MarkovZ, Zheleznykh, radhep} that rely on detection of radio Cherenkov emission via a process first proposed by Askaryan\\cite{Askaryan} and recently confirmed in the laboratory\\cite{Saltzberg}. The latter set includes detection of radio pulses generated in the lunar regolith\\cite{GorhamLN,GorhamLNSW}, geological salt deposits\\cite{wipp}, and Antarctic ice as observed {\\it in situ}\\cite{FrichterRM} or from a high altitude balloon\\cite{ANITA}. We report here on the first results from the Radio Ice Cherenkov Experiment (RICE), an array of dipole antennas located in-ice at the South Pole.  The concept behind RICE is illustrated in Figure \\ref{fig:concept}: a neutrino-nucleon interaction results in an electromagnetic cascade, which in turn produces a few nanosecond radio pulse with power concentrated around the Cherenkov angle. Detection of that pulse by several receivers with fast electronics allows for reconstruction of the interaction vertex and the energy/direction of the incident neutrino. Development and status of the RICE experimental configuration is summarized in Section \\ref{sec:overview}. \\begin{figure}[htpb] \\centerline{\\includegraphics[width=10cm]{evdisp2.eps}} \\caption{ RICE concept.  The RICE array (viewed from below) is drawn schematically to scale, showing cables and antenna locations. The small rectangle on the surface is the Martin A. Pomerantz Observatory (MAPO) which houses the RICE electronics. An electron neutrino is incident from above the horizon and interacts in the ice. An electromagnetic shower (hidden) results, which in turn generates an outgoing cone of radio Cherenkov radiation. Two cones, drawn at $\\theta_c \\pm 3^\\circ$, depict an approximate 3 dB range of signal strength. } \\label{fig:concept} \\end{figure}  The experiment requires accurate simulation of the expected signal, full calibration of an instrument with many interacting components, and development of techniques for event reconstruction and elimination of spurious backgrounds. Simulation and calibration are discussed in Section \\ref{sec:sigpath}. Additional information concerning our use of numerical techniques to model the Cherenkov signal may be found in Refs.~\\cite{RazzaqueSBMRS,RazzaqueSBMRSa}. Calibration of the RICE experiment is described in greater detail in Ref.~\\cite{Kravchenko}.  Event reconstruction (Section \\ref{sec:recon}) is important not only for the eventual analysis of neutrino observations, but also as a tool for separation of neutrino events from backgrounds. Background events result from either correlated or uncorrelated receiver hits. The former are due to impulsive transients associated with human activity in and around South Pole station\\cite{Frichter99}. Uncorrelated receiver hits result from thermal noise in the ice. The presence of continuous wave transmissions is readily detected, and is accounted for in the data analysis. Related issues are discussed in Section \\ref{sec:back}.  In development since 1995, the bulk of RICE operations has been devoted to calibration, understanding backgrounds, and developing strategies for dealing with those backgrounds. As a result, the experimental configuration has been dynamic - a situation which complicates both data-taking and analysis. We therefore chose to ``freeze\" the configuration for one month (August 2000), take data, and analyze the results. The analysis of August 2000 data is presented in Section~\\ref{sec:augytwok}. Refinement of the techniques used in the present analysis has been an ongoing process. As such, the present analysis is more complete than the preliminary analysis presented in Ref.~\\cite{Seckel01} and our present limit is somewhat stronger, although based on the same data.  Limited to a single month of data, the exposure analysed here is too short to provide a significant likelihood of finding actual neutrino events, at least given theoretical predictions. Our current search yields a null result, and we use it to place upper limits on proposed theoretical models. Accordingly, we have taken a conservative approach throughout our analysis, generally choosing simpler options when confronted with choices in the analysis, often resulting in a less restrictive limit. We present a broad survey of systematic uncertainties in Section~\\ref{sec:disc}.  RICE was originally conceived as an experiment to search for electron neutrinos via charged current interactions and the subsequent electromagnetic showers. Since these are reasonably well studied, the signal predictions are fairly robust and these events remain the primary subject of the current paper. It is understood, however, that at energies well in excess of 1 PeV, LPM effects reduce the efficiency for detecting electromagnetic showers. As a result, hadronic recoil showers, which are produced by neutrinos of all flavors in both charged and neutral currents, may be important for detection of neutrinos at the highest energies. Section~\\ref{sec:disc} includes a discussion of the efficiency for RICE to detect neutrinos via hadronic showers.  Finally, we summarize our discussion in Section~\\ref{sec:summary} and consider possible future applications of the radio technique at the South Pole. ",
        "conclusions": "\\label{sec:disc}  The search for ultra high energy neutrinos described in this paper yields zero signal candidates. Given the prototype nature of the current RICE deployment and the limited exposure used for the present analysis, this is not an entirely unexpected result. Accordingly, our discussion is focussed on the uncertainties that go into our analysis and on relaxing the limitations we imposed on the event types and geometry we consider.  \\subsection{Uncertainties} \\label{sec:uncertainties}  Potential systematic uncertainties are summarized in Table \\ref{tab:sys_sigma}. The Table is organized along the same lines as the discussion in Section \\ref{sec:sigpath}, where a detailed description of input to our Monte Carlo is described. Quoted uncertainties refer to particular quantities, not how they propagate through to an uncertainty in limits on any specific flux model.  Uncertainties in neutrino nucleon cross-sections reflect differences among available standard model perturbative QCD calculations. Similarly, uncertainties in shower simulations and signal generation reflect differences between GEANT 4 and ZHS simulations. A discussion of LPM issues is given below and leads into the discussion of hadronic showers and other event types. A complete understanding of attenuation in polar ice is still lacking. The rough factor of two quoted in the table is based on the considerations discussed in Section \\ref{sec:iceprop}. The 3 dB uncertainty in channel response is taken from Ref.~\\cite{Kravchenko}, as are estimates of timing resolution and our ability to reconstruct vertex location. Uncertainties for the on-line filter and off-line analysis efficiencies are based on Monte Carlo studies.  \\begin{table} \\begin{center} \\begin{tabular}{|c|c|c|}  \\hline Group & Item & $\\pm$ \\\\  \\hline Event Generator (Sec. \\ref{sec:showerrate}) & $\\sigma_{\\nu N}$ & $20\\%$  \\\\  & $(1-y)$ & $< 5\\%$ \\\\  \\hline EM pulse & shower simulation & $10\\%$\\footnote{Pulse amplitude}  \\\\ (Sec. \\ref{sec:empulse}) & spectrum & $10\\%$\\footnote{For RICE bandpass of 200-500 MHz} \\\\ & extrapolation: TeV $\\rightarrow$ PeV & $5\\%$ \\\\ & LPM effects & See text. \\\\  \\hline Propagation & attenuation $\\nu < 1 $ GHz & +0.3/-0.15 dB/100 m\\footnote{This corresponds to a factor of two uncertainty in $\\lambda_{att}$.}\\\\ (Sec. \\ref{sec:iceprop}) & attenuation $\\nu > 1$ GHz & \\footnote{Not critical given bandwidth limitations of antennas, cables, amps, and DAQ elements.} \\\\  & index of refraction & $< 1\\%$ \\\\  & reflections & small \\\\   \\hline Antenna and DAQ & overall channel response & 3 dB\\\\ (Sec. \\ref{sec:atot}) & overall channel timing & 25/50 ns\\footnote{25 ns for neutrino events. 50 ns for surface generated noise.} \\\\  & on-line filter efficiency & $5\\%$  \\\\  & dead time & $<5\\%$ \\\\  \\hline Analysis & $V_{eff}$ & 8\\%\\footnote{Statistical: the Monte Carlo runs produce roughly 150 detections in each $E_s$ bin.} \\\\ (Secs. \\ref{sec:veff}, \\ref{sec:recon}) & nearby vertex reconstruction & 10 m\\footnote{w/ assumed 50 ns pulse timing uncertainty, appropriate for noise rejection.}\\\\  & far vertex reconstruction & 0.1 R\\footnote{w/ assumed 50 ns pulse timing uncertainty, appropriate for noise rejection.}\\\\  & efficiency of analysis chain & $5\\%$  \\\\  \\hline \\end{tabular} \\caption{\\label{tab:sys_sigma} Estimates of systematic effects for different components of the RICE analysis.} \\end{center} \\end{table}  The largest uncertainties concern the attenuation in ice at frequencies of order 300 MHz and the channel-to-channel calibration of the receivers, as summarized by the gray bands in Figure \\ref{fig:eff_volume}. Uncertainties in calibration of signal strength or receiver sensitivity directly affect the threshold for RICE, but do not greatly alter the effective volume for shower energies above 1 EeV. This contrasts with the effects of modifying attenuation lengths, which would mostly affect sensitivity at high energies. These effects are illustrated by Table \\ref{tab:sys_event}. The Stecker and Salamon AGN model (a), for example, produces events at lower energy, and is seen to be most sensitive to changes in signal strength and sensitivity, whereas the topological defect (d) and GZK (e) models extend to higher energy and are more sensitive to uncertainty in $\\lambda_{att}$. \\begin{table} \\begin{center} \\begin{tabular}{c|c|cc|cc} Model & Nom & 2S & 0.5S & $2\\lambda_{att}$ & $0.5\\lambda_{att}$ \\\\ \\hline (a) & 0.0120 & 0.0433 & 0.0023 & 0.0220 & 0.0056 \\\\ (b) & 0.1384 & 0.2597 & 0.0572 & 0.3062 & 0.0526 \\\\ (c) & 0.0050 & 0.0079 & 0.0025 & 0.0118 & 0.0017 \\\\ (d) & 0.0182 & 0.0239 & 0.0123 & 0.0494 & 0.0059 \\\\ (e) & 0.0015 & 0.0021 & 0.0009 & 0.0037 & 0.0005 \\\\ \\end{tabular} \\caption{\\label{tab:sys_event} Expected number of events during a 333.3 hr exposure for different models and assumptions about systematic uncertainties. Models are labeled as in Fig. \\ref{fig:upper_limits}. ``Nom\" refers to the nominal case used for generating Fig. \\ref{fig:upper_limits}. ``2S\" and ``0.5S\" reflect the change in event rate if the signal strength were doubled or halved. ``2$\\lambda_{att}$\" and ``0.5$\\lambda_{att}$\" show the result due to doubling or halving the attenuation length.} \\end{center} \\end{table}  To account for the LPM effect, the RICE Monte Carlo uses the parameterization of Ref.~\\cite{AlvarezZ97}, which describes a narrowing of the Cherenkov cone due to a lengthening of the shower. The radiation pattern is based on an average shower profile, whereas fluctuations in longitudinal development could lead to ``non-gaussian\" tails for the angular pattern, which would enhance the effective volume. In addition, the parameterization is developed in the Fraunhoffer regime, valid for $R \\gg l^2/\\lambda$, where $R$ is the distance to the observer, $l$ is the shower length, and $\\lambda$ is the wavelength of the radiation. As $l$ increases with the LPM effect, a more accurate treatment will lead to a broader radiation pattern\\cite{BuniyR,AlvarezVZ}, which may enhance the effective volume for high energy showers.  \\subsection{Additional event types and hadronic showers} \\label{sec:hadronic}  We have restricted our main analysis to electromagnetic showers produced in charged current events of electron neutrinos. The reason for this is two fold: a) simplicity, and b) this process is better studied in the literature. Even so, expanding the analysis to include neutral currents and other neutrino flavors may be expected to increase the rate of neutrino interaction candidates by a factor of 4.5. There is a 50\\% increase due to including neutral currents and a factor of 3 for including all flavors. Additional increases may occur if we include upward neutrinos from just below the horizon and extra efficiency for $\\nu_\\tau$ due to regeneration in propagation through the Earth and $\\tau$ decay in the ice near the RICE array. Table \\ref{tab:evtype} summarizes restrictions we have placed on our analysis, and the potential effect on event rates if the restriction were removed.  \\begin{table} \\begin{center} \\begin{tabular}{|c|c|}  \\hline Restriction & Effect of relaxation\\footnote{Estimates are made at $E_{\\nu_e} = 1$ EeV, unless stated, where signal from hadronic shower and EM shower are comparable (see Figure \\ref{fig:veff_lpm}).} \\\\  \\hline $\\theta_z > 90 \\deg$ & Increase rates 10\\%\\footnote{The Earth is not totally opaque just below the horizon, so rates increase slightly, even at high energies.} \\\\  no hadronic shower in $\\nu_e$ charged current events & 0-100\\%\\footnote{depends strongly on energy and geometry.}\\\\  CC events only & + 50\\% per flavor\\\\  $\\nu_e$ events only & $\\times$ 3 \\\\  $\\tau$ decay not included & increase by 1 charged current channel for $E < 20$ PeV\\\\  $\\tau$ regeneration not included & increased event rate for upward neutrinos\\\\  \\hline \\end{tabular} \\caption{\\label{tab:evtype} Summary of event type and geometry considerations in RICE analysis.} \\end{center} \\end{table}   As described in Section \\ref{sec:showerrate} the hadronic energy is typically a quarter of that available on the leptonic side of an event. Most of the energy ends up as electromagnetic energy, primarily through $\\pi^0$ decay; however, that process is not effective until the average pion energy drops below $\\sim 6$ PeV, similar to the energy where LPM effects start to have a significant effect on the longitudinal evolution of the shower. As a result, hadronic showers, although decreased and more variable in energy, do not suffer the decrease in solid angle acceptance for electromagnetic showers related to the LPM effect.  Given the potential importance of hadronic showers, we have recalculated the RICE effective volume under the assumption that LPM effects may be ignored - i.e. the signal strength scales linearly with shower energy and the angular pattern is energy independent. The result is shown in Figure \\ref{fig:veff_lpm}. From the Figure, one may infer that the hadronic shower provides a stronger radio Cherenkov signal than the electromagnetic shower for neutrino energies above $\\sim 1$ EeV, for typical charged current $\\nu_e$ events. \\begin{figure}[htpb] \\begin{center} \\includegraphics[width=12cm]{veff_lpm.eps} \\end{center} \\caption{ Effective volume with and without the LPM effect. The solid curve is the same nominal result with LPM shown in Figure \\ref{fig:eff_volume}. The dotted curve shows $V_{eff}$ calculated without the LPM effect. The dashed curve shows the non-LPM result, but shifted in energy by a factor of 4. This roughly mimics the relative response of RICE to hadronic showers (as compared to electromagnetic showers from the same energy neutrino), where the average inelasticity is $<y> = 0.2$. } \\label{fig:veff_lpm} \\end{figure}  It is now straight forward to replicate the main analysis, with appropriate modifications, to study the potential improvement in sensitivity to be gained by considering hadronic showers produced in charged and neutral currents of all neutrino flavors. These modifications include production of hadronic showers with energy $E_s = y E_\\nu$ as outlined in Section \\ref{sec:showerrate}, and use of the ``no-LPM\" $V_{eff}$ shown as the dotted curve in Figure \\ref{fig:veff_lpm}.  There are three nominal choices to make for the neutrino flavor content, which all lead to the same event rates. If there is no mixing, the flavor content is determined by the physics of $\\pi$-decay, which yields the well known relation $\\phi_{\\nu_\\mu} \\approx 2 \\phi_{\\nu_e}$. If there is mixing amongst $\\nu_\\mu$ and $\\nu_\\tau$ as suggested by the SuperKamiokande atmospheric neutrino analysis\\cite{superk-atmospheric}, but not $\\nu_e$, then the event rate of hadronic showers is unchanged, since $\\nu_\\mu$ and $\\nu_\\tau$ produce the same hadronic showers at these energies. If there is full mixing of all flavors, then the rate of EM showers is unchanged since the $2:1$ flavor ratio ($\\nu_\\mu+\\nu_\\tau : \\nu_e$) at production is unchanged by mixing. For any of these scenarios, Table \\ref{tab:had_event} gives event rates including hadronic events. For the lower energy AGN models, including hadronic events increases the event rate by factors of 2-4. For the GZK (e) and topological defect (d) models, the increase is significantly larger - 7 and 16 respectively.  We have not modeled $\\tau$ decay events. To be observed, the $\\tau$ would have to decay within a km or so of its production site, corresponding to $E_\\tau \\approx 20$~PeV. Thus, charged current $\\nu_\\tau$ events presumably contribute an event rate comparable to the other terms in the Table, but which is suppressed at high energy by the probability $P \\sim \\frac{20 {\\rm PeV}}{E+20 {\\rm PeV}}$ for the $\\tau$ to decay in the ice. A proper estimate of this term requires a separate Monte Carlo to deal with the stochastic nature of $\\tau$ decay. Also, since we confine our analysis to the downward flux of neutrinos, we do not allow for the detection of an upward flux of regenerated $\\nu_\\tau$\\cite{SaltzbergH}, or $\\nu_e$ and $\\nu_\\mu$ produced in $\\tau$ decay\\cite{BeacomCK}. Except near the horizon, these are expected to be of fairly low energy compared to the main RICE sensitivity.  \\begin{table} \\begin{center} \\begin{tabular}{c|c|cc|c} Model & $\\nu_e$ EM & $\\nu_e$ H & $\\nu_\\mu + \\nu_\\tau$ H & Total \\\\ \\hline (a) & 0.0120 & 0.0049 & 0.0098 & 0.0268 \\\\ (b) & 0.1384 & 0.1096 & 0.2193 & 0.4673 \\\\ (c) & 0.0050 & 0.0063 & 0.0126 & 0.0238 \\\\ (d) & 0.0182 & 0.0962 & 0.1923 & 0.3067 \\\\ (e) & 0.0015 & 0.0029 & 0.0057 & 0.0100 \\\\ \\end{tabular} \\caption{\\label{tab:had_event} Expected number of events during a 333.3 hr exposure for different models, including hadronic showers. Models are labeled as in Fig. \\ref{fig:upper_limits}. ``$\\nu_e$ EM\" is the nominal event rate shown in Fig. \\ref{fig:upper_limits} for our main analysis. ``$\\nu_e$ H\" gives the event rate due to the hadronic recoil showers from $\\nu_e$ charged current and neutral current events. Similarly, ``$\\nu_\\mu + \\nu_\\tau$ H\" gives the event rate due to $\\mu$ and $\\tau$ neutrinos. For an isoflavor model this is just twice the $\\nu_e$ H rate. ``Total\" gives the sum of the three contributions, as appropriate for either an isoflavor model or a conventional source model with no mixing and $\\phi_{\\nu_\\mu} \\approx 2 \\phi_{\\nu_e}$. } \\end{center} \\end{table}      Using data from the RICE antenna array collected during August 2000, we have searched for nanosecond radio pulses resulting from neutrino interactions in ice. We find no such events, which provides the basis for the diffuse flux limits on ultra-high energy electron neutrinos shown in Figure \\ref{fig:upper_limits} and Figure \\ref{fig:uplim2}. This is the primary result of this paper.  At the same time, we have demonstrated the efficacy of techniques designed to eliminate anthropogenic backgrounds and thermal noise.  There are many ways by which either a) the limits we have derived can be improved to the level where they seriously constrain models, or b) the experiment results in actual UHE neutrino detection. The most obvious of these is to increase the exposure and validate sensitivity to hadronic showers. Archived data presents a potential increase of more than an order of magnitude increase in exposure. Based on the estimates in Table \\ref{tab:had_event}, including hadronic cascades would increase the event rates for typical AGN models by factors of 2-4, and by factors as high as 10-15 for GZK and topological defect models. Combined, these factors may achieve a factor of nearly 100 in sensitivity. This would severely test several models - especially the neutrino fluxes associated with top-down models proposed to explain the purported excess of cosmic rays with energies above the GZK cutoff. It would, however, remain about a decade short of testing the Waxman-Bahcall bound\\cite{WaxmanB} or probing the ``guaranteed\" flux of neutrinos produced in the GZK process.  To further improve the limits requires improving the acceptance of the experiment by reducing the threshold for event detection, or by increasing the number of antennas and the area over which they are deployed. The main impact of lowering the threshold is to increase the sensitivity at modest energies, 1-100 PeV, but, as illustrated in Figure \\ref{fig:eff_volume}, the gains at higher energy are quite modest. As signal attenuation determines the ``horizon\" for a single antenna, to improve the RICE sensitivity for energies in excess of an EeV requires deployment of antennas over a wider area.  There are a number of technological improvements that could dramatically improve the sensitivity for energies below 100 PeV, including making use of correlated cross-polarized antennas, increasing the bandwidth, making use of matched filters to maximize signal to noise, and deploying larger numbers of antennas to increase the likelihood that the Cherenkov cone from a true neutrino would intersect four or more antennas.  To ensure that the experiment remains signal limited (as opposed to background limited) the current noise event rejection must be improved without significant loss of efficiency for neutrino events. This should be achievable by simple strategies. For example, increasing the number of antenna hits required in coincidence from 4 to 5, results in more than a hundredfold reduction in thermal noise events, but only a $\\sim 15\\%$ reduction in neutrino induced events. Similarly, a modest increase in the restrictions of the on-line veto or the analysis cuts on event reconstruction, would reduce impulsive backgrounds significantly without significant reduction in neutrino detection efficiency.  Judging by the summary in Table \\ref{tab:sys_sigma}, uncertainties associated with RICE need to be reduced as well. Most important is the need for a better determination of the radio attenuation length in ice. We believe the current treatment of $\\lambda$ is on the conservative side by some 30\\%-50\\%, which would potentially increase $V_{eff}$ by a factor $\\sim 2$ for $E_\\nu > 1$ EeV. Similarly, we believe the parameterization of the LPM effect taken from Ref. \\cite{AlvarezZ97} may lead to an underestimate of $V_{eff}$.  Finally, we look forward to the time when true neutrino events are detected by RICE, at which point critical issues will turn from acceptance and noise rejection to the ability to reconstruct events and identify sources and neutrino flavor on an event by event basis. To that end we are improving the timing analysis needed for event reconstruction. Improving the amplitude calibration is critical here as well, since signal strength plays a role in reconstruction of the Cherenkov cone.  To summarize, we believe the outlook for RICE, and the radio Cherenkov technique in general, is quite bright. Analysis of data already in hand should significantly improve present limits on $\\nu_e$ fluxes. Inclusion of other flavors and interaction modes will further constrain models. We expect the experiment to remain signal limited, not background limited, at least down to levels where even conservative flux models may be observed. The ability to reconstruct the vertex and incident neutrino direction should permit searches for point sources. Similarly, arrival times are sufficiently resolved to conduct coincidence studies between RICE events and gamma-ray bursts. Finally, we expect that the technique should be extendable to place new limits on ``light\" relativistic magnetic monopoles\\cite{Wick}, as well as neutrino interactions, including production of micro black holes\\cite{AlvarezFHHH,AnchordoquiFGS,KowalskiRT,JainKMPR}."
    },
    "0206/astro-ph0206001_arXiv.txt": {
        "abstract": "Having been operational at Kitt Peak for more than a year, the prototype of the Hungarian Automated Telescope (HAT-1) has been used for all-sky variability search of the northern hemisphere. The small autonomous observatory is recording brightness of stars in the range of $\\rm I_c\\approx 6\\--13^m$ with a telephoto lens and its $9\\arcdeg\\times9\\arcdeg$ field of view (FOV), yielding a data rate of $\\sim10^6$ photometric measurements per night. We give brief hardware and software description of the system, controlled by a single PC running RealTime Linux operating system (OS). We overview site-specific details, and quantify the astrometric and photometric capabilities of HAT. As a demonstration of system performance we give a sample of 60 short period variables in a single selected field, all bright, with $\\rm I < 13^m$, of which only 14 were known before. Depending on the observing strategy, search for extrasolar planet transits is also a feasible observing program. We conclude with a short discussion on future directions. Further information can be found at the HAT homepage: {\\tt http://www-cfa.harvard.edu/\\~{ }gbakos/HAT/}. ",
        "introduction": "The idea of automating ground based observational astronomy goes back more than two decades. Minimizing manpower can assure uniform and massive data flow with low budget and the absence of human mistakes.   The increasing number of robotic telescopes\\footnote{ See eg.~http://alpha.uni-sw.gwdg.de/\\~{ }hessman/MONET} (capable of computer controlled multiple observations) have been used for a broad range of projects, such as astrometry: Carlsberg Meridian Telescope \\citep{carlsberg}; photoelectric photometry of pre-selected targets: Fairborn Observatory \\citep{Fairborn}; supernova search: KAIT \\citep{KAIT}; GRB follow-up: ROTSE \\citep{Akerlof00,ROTSEIII} and LOTIS \\citep{Lotis,SuperLotis}, exo-planet searches: STARE \\citep{STARE}, Vulcan \\citep{Vulcan}; and asteroid searches: TAOS \\citep{TAOS}. The variety of targets is usually narrow, looking only for specific timescales, light-curve shapes and intensity ranges.  Although most projects concentrating on special targets gain a huge amount of photometric data, only few of them are capable of presenting their by-products to the astronomical community, e.g.~OGLE \\citep[e.g.~][]{Wozniak02}, MACHO \\citep{Allsman01} and ROTSE \\citep{Akerlof00}.  Initiated by ideas of \\citet{BP97}, the All-Sky Automated Survey's \\citep[ASAS;][]{GP97} approach is different, in that the final goal has been photometric monitoring of {\\em all bright stars} in a major part of the southern sky down to $\\rm I\\approx 14^{m}$. Using a fully automated but inexpensive system consisting of an amateur-class CCD, a small telephoto lens and an equatorial mount, ASAS presented catalogues of 4000 bright variables from a $300 \\sq\\arcdeg$ area of the southern sky, {\\em 96\\% being new discoveries} \\citep{GP98,GP00}. The upgraded ASAS-3 will produce an order of magnitude increase in the data flow.\\footnote{http://www.astrouw.edu.pl/\\~{ }gp/asas/asas\\_asas3.html} The incompleteness of our knowledge on bright variable stars was reinforced by \\citet{Akerlof00} who discovered 1781 new variables in a $2000\\sq\\arcdeg$ area.  Why is general variability study of bright objects important? Several answers can be found in \\citet{BP97,BP00}, and others can be added. Variable stars are essential for testing stellar structure and evolution theories, examining galactic structure or establishing the extragalactic distance scale. Only bright variables are within the range of high resolution spectroscopy, parallax and proper motion measurements. Our knowledge of issues related to variable stars (e.g.~distance scale) can be refined by the combination of detailed study of close-by, bright objects and of equidistant, homogeneous samples (e.g.~OGLE \\--- Galactic bulge, LMC, SMC). Serious incompleteness at the bright end affects all conclusions. A systematic, well-calibrated survey presents clean, statistically valuable samples with well defined limits for different subtypes of variable objects. A reliable database with sufficient and ever-growing time-span of light-curves can be used as an archive, for e.g., correlating optical variability with X-ray observations, made by satellites. It can be a valuable input to schedule big-telescope and space-mission observations, where telescope time is limited, or prior and longer-term data on field variables is necessary (e.g.~the Kepler mission), furthermore, it can provide them with a real-time alert system of rare events. Such events can be nova explosions, helium flash of a star \\citep[Sakurai's object:][]{sakurai1,sakurai2}, super-outbursts of dwarf-novae \\citep[WZ Sge:][]{wzsge1}.  To mention specific examples, observational data is scarce for spotted red subgiant variables (RS Cvn, FK Com), which are crucial in understanding the stellar magnetic cycles. Detached eclipsing binaries (through their stellar mass, radius and luminosity determination) can be perfect distance and age indicators, if nearby systems are properly calibrated. Samples of such objects in the solar neighborhood are sparse \\citep{BP97}, partly due to their short and narrow eclipses, and lack of observational data (see Fig.~\\ref{fig:twovar} for our light-curve of a {\\em semi}-detached binary). Bright contact binaries exhibit similar incompleteness, although long-term observations could reveal interesting phenomena, such as the transition to semi-detached state.  The Hipparcos Space Astrometry Mission presented us with discovery of a few thousand new bright variables, and the Hertzsprung-Russell (HR) diagram was described in terms of luminosity stability at the millimagnitude level \\citep{hipvarcmd}. However, Hipparcos observed only selected stars (120000 or $3/\\sq\\arcdeg$), and the variable star sample is further limited by the $\\sim 110$ epochs per star on average and cut-off at $\\rm I\\approx9^m$.  Mapping the location of the large variety of pulsating variables on the HR diagram is still far from being complete. Addressing phenomena, such as long-term period and amplitude modulations (e.g.~Blazhko effect of RR Lyrae), evolution of the pulsational status of a star, is possible only by long-term and {\\em homogeneous} observations. Given the huge data-flow, interesting phenomena are expected to emerge, for instance further observational evidence for chaos in W Vir and RV Tau stars \\citep{Buchler95}, triple-mode variables \\citep[GSC 40181807:][]{triplegsc}, Cepheids which stop pulsating \\citep[V19 in M33:][]{Macri01} or strong amplitude modulation of Cepheids \\citep[V473 Lyr:][]{Burki86}. Long-term monitoring of semi-regular and Mira variables is needed to disentangle multiperiodicity and systematic amplitude variations \\citep[e.g.~][]{kiss00}. The sample of the recently established $\\gamma$ Dor subtype (oscillations in non-radial gravity-mode) consists of only a few dozen stars. One example of the possibilities is our HAT-1 light-curve of the triple-mode pulsator AC And (Fig.~\\ref{fig:twovar}).  \\begin{figure}[!h] \\epsscale{1.0} \\plotone{twovar.eps} \\caption{Phased light-curves of Algol-type, semi-detached eclipsing binary TT And, and triple-mode pulsator AC And from HAT observations (See \\S \\ref{sec:photprec}). ``Chaotic'' appearance of the AC light-curve is due to its triple-mode behavior. Phase was computed from its longest period. \\label{fig:twovar}} \\end{figure}  While the number of robotic telescopes  is around a hundred, there are only a few completely autonomous observatories, where the human supervision is eliminated, and all auxiliary appliances (dome, weather station, etc.) are under a reliable computer control. These observatories can be installed to remote sites with adequate astro-climate and infrastructure (electricity, Internet) without need of an on-site observer, and daily maintenance. The station can be monitored via Internet, and operation is not a bottleneck any more.  HAT is such a small autonomous observatory intended to carry out a northern counterpart of ASAS, i.e., a variability study of the northern sky. HAT was developed and constructed by one professional and three amateur astronomers\\footnote{ G.~Bakos (astronomical considerations and software), J.~L\\'az\\'ar (software development; www.xperts.hu), I.~Papp (electronic design) and P.~S\\'ari (mechanical engineering)} in Hungary, and has been fully operational since May, 2001 at Steward Observatory, Kitt Peak, Arizona. HAT is controlled by a single Linux PC without human supervision.  The 180mm focal length and 65mm aperture of the telephoto lens, and the $\\rm 2K\\times2K$ CCD yields a wide FOV: $9\\arcdeg\\times9\\arcdeg$ on the sky. Our typical exposure times allow us to study variability of the $\\sim20000$ objects per field brighter than $\\rm I_c\\approx13^m$ with few percent precision ($\\rm 0.01^m\\--0.05^m$), and few minute to one year time-resolution. As an extreme of the possible observing tactics, HAT is capable of recording the brightness of {\\em every locally and seasonally} visible star in the range of $\\rm I_c\\approx 6\\--13^m$ in every second day. Our limiting magnitude and photometric precision (See \\S\\ref{sec:photprec}) corresponds to the following detection {\\em cut-offs} for a few selected variability types (ranges indicate that the distance limit depends on the luminosity within the type): $\\gamma$ Doradus stars ($\\rm \\langle700pc\\rangle$), $\\delta$ Scuti ($\\rm 1\\--2.5kpc$), RR Lyrae ($\\rm 3kpc$), Cepheids ($\\rm 10\\--150kpc$), Miras and semiregular variables ($\\rm \\langle60kpc\\rangle$).  The limits are based on the luminosity of the sources, and \\--- especially at the larger values \\--- overestimate true detection cut-offs, as they do not take into account reddening and crowding.  With the above specifications, HAT is also suitable for exo-planet search via transits, which is also included in our program. However, we concentrate on a broader range of variabilities: a large fraction of the sky is monitored sparsely, and few selected fields frequently so as to have preliminary results on short-period changes.  We expect that our survey will contribute to most of the aforementioned issues related to variability search (public archive, alert system), and supplement the incomplete variable classes by new discoveries.  Our current data rate is roughly $10^6$ photometric measurement, or two Gigabytes of raw data per night. Typically a few percent of the sources are variable, i.e., variable light-curves are supplemented by $\\sim 20000$ data points each night.  HAT is monitoring only a fraction of the northern sky, but given the fact that it is an off-the-shelf system there is perspective for installation of new units in the near future.  The paper is arranged as follows: \\S \\ref{sec:hardware} gives an overview on the hardware, \\S \\ref{sec:software} describes our software environment, \\S \\ref{sec:astro} quantifies the pointing precision of HAT, \\S \\ref{sec:installation} and \\S \\ref{sec:obs} summarize our site-specific installation at Kitt Peak and observations in the past one year, \\S \\ref{sec:photprec} estimates our photometric precision, \\S \\ref{sec:summary} gives summary and future directions.  \\placefigure{fig:twovar} \\notetoeditor{It would be important to place Fig.~\\ref{fig:twovar} in the introduction part, so the reader sees some interesting results before struggling through the techincal part} ",
        "conclusions": "\\label{sec:summary} We described a small, autonomous observatory, which has been working for one year at Steward Observatory, Kitt Peak. In spite of the small telephoto lens used as the ``telescope'', it can perform massive photometry of bright sources. It completed on the order of $20000$ pointings to different objects, yielding a data flow of $\\sim 10^6$ photometric measurements per night, and proved to work reliably. Using 5\\% of our data, and a single selected field, a few dozen bright ($I<13$) variables were found, further reinforcing the incompleteness even on the bright end of previous variability searches. Some parameters of the system are summarized in Table~\\ref{tab:spec}.  HAT-1 is capable of monitoring a {\\em small fraction} of the sky with sufficient time resolution. Multiple filters would not only ease classification of sources based upon their colors, but due to differential refraction (vs.~color), the photometric precision would be also improved.  Wider aperture would increase the incoming flux, and either our time-resolution at constant photometric precision would drop, or our limiting magnitude with the same exposure times would be expanded. This would not be necesseraly {\\em improvement} of the system, but in some sense the target of observations would be different. HAT is designed to operate in a network. It is an off-the-shelf system of $\\rm \\sim15K$ USD cost, plus the CCD. These make multiple installations easy. It is also a flexible, multi-purpose observatory, not only eligible for all-sky monitoring, but could be used as a photometric monitor station at bigger observatories.  Current bottleneck is neither hardware development or operating HAT, but rather efficient data reduction, archiving and web-availability of the data. After efficient software can handle the data flow and overcome the difficulties, we plan to upgrade HAT with more filters, bigger telescopes, and to install multiple stations.  \\begin{table}[h] \\caption{Specification of HAT-1 at Kitt Peak. \\label{tab:spec}} \\begin{tabular}{lr}\\hline \\multicolumn{2}{c}{Equatorial mount}\\\\\\hline\\hline Max.~RA slew & $\\rm 2\\arcdeg/$sec\\\\ Max.~Dec slew & $\\rm 5\\arcdeg/$sec\\\\ RA resolution & $\\rm 1\\arcsec/$step\\\\ Dec resolution & $\\rm 5\\arcsec/$step\\\\ Max.~tel.~diam. & 20cm\\\\\\hline \\multicolumn{2}{c}{CCD (Apogee AP10)}\\\\\\hline\\hline Dimensions & $\\rm 2K\\times2K$, $\\rm 14$ micron\\\\ Gain & 10 e-/ADU\\\\ Readout noise & 2 ADU\\\\ Dark current\t&\t$\\rm 0.05ADU/sec$ at $-15\\C$\\\\ Readout time & $\\rm <10s$\\\\ Cooling & Peltier ($\\rm \\Delta T=32\\C$)\\\\\\hline \\multicolumn{2}{c}{Lens (Nikon 180mm f/2.8)}\\\\\\hline\\hline Aperture & 65mm\\\\ Plate scale & $\\rm 16\\arcsec/pixel$\\\\ Field of view & $9\\arcdeg\\times9\\arcdeg$\\\\\\hline \\multicolumn{2}{c}{Pointing precision}\\\\\\hline\\hline Homing  & $<0.5\\arcmin$\\\\ Tracking & $\\rm <0.5sec/2hr$\\\\ Absolute & RA:$\\rm 20sec$, Dec: $1\\arcmin$\\\\ Repeatability & RA: $\\rm 7sec$, Dec: $1\\arcmin$\\\\\\hline \\multicolumn{2}{c}{Photometry ($\\rm I_c$ band, 8min time res.)}\\\\\\hline\\hline Accuracy (absolute) & $\\rm \\approx 0.05^m$\\\\ Precision $\\rm (I_c < 10^m)$& $\\rm \\lesssim 0.01^m$\\\\ $\\rm I_c=11^m$ & $\\rm 0.02^m$\\\\ $\\rm I_c=12^m$ & $\\rm 0.05^m$\\\\\\hline\\hline \\end{tabular} \\end{table}"
    },
    "0206/astro-ph0206237_arXiv.txt": {
        "abstract": "In this paper the potential of high resolution spectroscopy of nearby AGN with XEUS is discussed.  The focus is upon the energy resolution that is needed in order to disentangle the different spectral components.  It is shown that there is an urgent need for high spectral resolution, and that a spectral resolution of 1~eV, if possible, leads to a significant increase in diagnostic power as compared to 2~eV resolution. ",
        "introduction": "Since the launch of Chandra and XMM-Newton, with their high resolution spectrometers, our insight into the astrophysics of Active Galactic Nuclei (AGN) has changed dramatically.  It is now possible to study in detail the geometry, dynamics and physical state of the warm absorber, as well as the underlying continuum spectrum, including the exciting possibility of relativistic emission lines.  While with XMM-Newton and Chandra high resolution, high signal-to-noise ratio spectra of the brightest (and in general the nearest) AGN can be taken, the large resolving power and effective area of XEUS will allow us to study AGN spectra out to large redshifts or low intrinsic luminosities.  It will also allow time-resolved spectroscopy of the most rapidly varying AGN.  A detailed understanding of the astrophysics and spectral signatures of the nearest AGN is an absolute requirement in order to understand the properties of the most distant AGN that will be observed by XEUS, which will have much noisier spectra. In this contribution  spectral simulations  are presented that show the potential and limitations of AGN spectroscopy with XEUS. ",
        "conclusions": ""
    },
    "0206/astro-ph0206015_arXiv.txt": {
        "abstract": "We present Very Long Baseline Array (VLBA) observations of the radio source J1625+4134 at 22 and 15 GHz and analyze them in concurrence with other existing VLBI data on this source. The high resolution images at 15 and 22 GHz show a short and bending jet which has about $270\\degr$ difference in position angle with the northern jet detected at lower frequencies. The new high resolution data, combined with the data available in the literature, allow us to estimate the spectral index of the components and identify one of the compact components as the VLBI core based on its flat spectrum between 5 and 22 GHz. Relative to this core component, the jet appears to be bi-directional. The proper motion measurement of the component C2 and the estimate of the Doppler boosting factor suggest that the orientation of the jet is close to the line of sight. The projection effect of an intrinsically sharply bending jet within a few mas from the core or the erratic change in the nozzle direction of the jet may account for the uncommon bi-directional structure of the jet in J1625+4134. ",
        "introduction": "The radio source J1625+4134 (1624+416; 4C\\,41.32) was identified as a quasar at the redshift of $z=2.55$ \\citep{Pear88, Hewitt93}. It has a $22^{m}$ optical magnitude, and has been detected at the near infrared K band \\citep {Lebo83}. The source shows a flat spectrum between 178 MHz and 37 GHz with the spectral index $\\alpha = -0.34$ ($S_\\nu\\propto\\nu^{\\alpha}$). At the VLA angular scale, the source extends to about $1\\arcsec$ in the northern direction with a very weak polarization indicative of the magnetic field parallel to the jet structure \\citep{Per82, ODea88}.  The quasar J1625+4134 has been observed with VLBI at various frequencies: 1.67 GHz \\citep{Pol95}, 2.32 and 8.55 GHz \\citep{Fey97}; 5 GHz \\citep{Pear88, Foma2000, Lis2001a}. On the VLBI scale, its structure is rather uncommon: a short strong jet towards the southwest and a long but weak jet emission extended to the north--northwest.  The VLBI images at 1.67 and 2.32 GHz are similar, the jet of J1625+4134 extends to the north to over 30 mas along the same position angle PA~=~$350 \\degr$ as that in the VLA images. Model fitting of the VLBI data revealed a component in the southwest region which was not clearly separated from the central compact emission in the images probably due to the limited angular resolution \\citep{Pol95, Fey97}. The VLBI images at 5 and 8.55 GHz showed a compact core, a short curved jet in the southwest region and a weak emission to the north. The high resolution VSOP image at 5 GHz showed the structure of the southwest jet only, while the image made with the ground-only array (VLBA) of the same VSOP observation  exhibited the two-sided jet structure \\citep{Lis2001a} extending both in southwest and north directions. This kind of sharp difference between VLBI structures at low frequencies (and also lower angular resolutions) and higher frequencies (and higher angular resolutions) is relatively rare in AGNs. Another similar case has been found in the high redshift quasar 1351-017 ($z=3.707$, Frey et al. 2002).  There are several possible explanations for this morphological difference seen in the jet structure of J1625+4134 at different frequencies and angular scales. It has been found that milliarcsecond-scale jets in many radio loud sources is misaligned with respect to the structure at the arcsecond scale \\citep{Pear88, Hong98, Cao2000}. The large $\\Delta$PA (the difference between the jet position angles measured on mas and arcsecond scales) may be due to the projection of a bending jet or a helical jet \\citep{Conway93}. If the brightest component in the images is the VLBI jet base, J1625+4134 can be classified as a quasar with a sharp bend. For the jet moving away from the core toward the northern emission, $\\Delta $PA is about $90\\degr$ \\citep{Lis2001b} or $270\\degr$ depending on the jet bending direction.  Alternatively, the brightest component in the 5 and 8.5 GHz images could be not the core of J1625+4134. Rather, the core could be weak and located, e.g., at the end of the southwest emission. In this case the brightest emission in the 5 and 8.5 GHz images may be caused by interaction of the jet with the ambient medium. The jet turns for about $90\\degr$ to the north in this region of strong jet-medium interaction. Such the interaction will produce a strong shock, which further enhances the emission and flattens the spectrum.  The third possibility is that the source has a twin-jet structure with the radio jet lying close to the sky plane. In this case, the relativistic boosting effect is not important, the jet and counter-jet are both detectable. The jet emission should be more symmetric at the higher frequencies due to the reduced free-free absorption of the obscuring accretion disk or torus.  In this paper we present and discuss the results of VLBA observations at 22 and 15 GHz. In Section 2, we describe the observations and data reduction. The results are presented in Section 3. In Section 4, we discuss the jet morphology. Throughout this paper we adopt the Hubble constant $H_{0}$ = 65 km\\, s$^{-1}$\\,Mpc$^{-1}$ and the deceleration parameter $q_{0}=0.5$, assuming the cosmological constant $\\Lambda = 0$. With these parameters, 1 milliarcsecond (mas) corresponds to 5.91~pc at the redshift of $z = 2.55$.   \\section {OBSERVATIONS AND DATA REDUCTION}  The observations of J1625+4134 at 22 GHz were conducted with the VLBA on 1 March 2000, the total on-source time was about 3 hours. The LCP signals were recorded in 4 IF-bands for a total bandwidth of 32 MHz with 2-bit sampling. Two strong calibrators, J1635+3805 and J1640+3946, were observed too. The data were then correlated at the VLBA correlator in Socorro (NM, USA) and calibrated and fringe-fitted using the NRAO Astronomical Image Processing System (AIPS) software. Initial amplitude calibration was carried out with the system temperature measurements and the NRAO-supplied gain curves, before the fring-fitting, the data are corrected for 2-bit sampling errors and atmospheric opacity. The imaging and self-calibration were carried out in DIFMAP package \\citep{Shep94}. The constant gain corrections of the stations during the self-calibration are well in agreement with those for the two calibrators.  The source was also observed as part of the VLBA survey at 15 GHz in support to the VSOP Survey Programme \\citep{Gur2002}. The observation was conducted on 2 January 1999 with a total bandwidth of 64 MHz using 8 IF-bands and 1-bit sampling; the total on-source time was about 35 minutes (7 scans $ \\times$ 5 minutes). The data on J1625+4134 were calibrated and fringe-fitted alongside with other survey sources, but for the purpose of this work the image was produced independently using AIPS.  \\section {RESULTS}  \\subsection {Images}  The left panel in Fig.~1 presents the naturally weighted image of J1625+4134 at 22~GHz. The estimated thermal noise ($1\\sigma$) is about 0.3 mJy/beam for 3 hours on-source integration, the lowest contour in the image is 0.9~mJy/beam. The core-jet morphology is consistent with those observed at 5 and 8.5~GHz, the jet bends smoothly to the southwest from the brightest spot and the emission becomes diffuse at about 3 mas.  The naturally weighted image at 15 GHz is shown in the right panel of Fig.~1. The estimated thermal noise ($1\\sigma$) is about 0.4~mJy/beam for 35 minutes on-source integration, the lowest contour in the image is 1.0~mJy/beam. The general source structure at this frequency is similar to that at 22 GHz. However, one can notice traces of a weak emission at 15 GHz north-northwest at a distance of about 3 mas from the brightest component coinciding with the northern components seen in the images at 5 and 8.5~GHz \\citep{Pear88, Lis2001a, Fey97}.  The 5-component model represents well the core-jet structure at 22 GHz and is listed in Table~1. The calibrated $u,v$-data at 15~GHz were also fitted by 5 Gaussian components (see Table~1), but there is a very low surface brightness emission in the northern region in the residual map.  The estimated flux density of this low brightness emission is about 14 mJy. In order to compare our results with other observations, we follow the labeling convention of the jet components introduced by Lister et al. (2001a) in their 5 GHz VSOP image. The component from our model called C2B and present in both 15 and 22~GHz data (see Fig.~1 and Table~1) is absent in the VSOP image at 5~GHz by Lister et al. (2001a). In Table~1, we also give the error estimations of the flux density and the position of the model components. The flux density errors of the model components include the calibration uncertainty as well as the spread in model-fitting results obtained in AIPS and DIFMAP. The error estimation of the component positions are determined by the differences between different model-fitting procedures and the uncertainties estimated by the formula introduced by Fomalont \\cite{Foma88}.  The component D is the strongest and the most compact at both 15 and 22~GHz. It has a flat spectrum at the higher observing frequencies, which, following the existing convention, allows us to identify this component as a 'core'. At the brightness level of 1~mJy/beam, no evidence of the counter-jet emission was found in the opposite direction of the bright jet. The weak emission in the north-northwest region at 15 GHz does not seem to resemble a collimated counter-jet, although such an interpretation cannot be ruled out at this stage.   \\subsection {The distribution of the VLBI component positions}  The right panel in Fig.~2 presents a plot of the position of the jet components measured at different frequencies in the quasar J1625+4134 relative to the core component D. The two-sided jet structure consists of a short bright jet bending from the west to the southwest at a few mas from the core and a long weak jet extending to the north for more than 25~mas. If the curved southwest jet continuously bends to the north, the overall bending is at least $\\Delta {\\rm PA} \\approx 270\\degr$. Such an extreme bending on the VLBI scale is relatively rare in the known radio structures in AGNs.  The left panel of Fig.~2 shows the enlarged southwest jet. Within the inner 2 mas from the core, the jet components at 8, 15 and 22 GHz follow the same trajectory moving away from the core. But there is about $15\\degr$ difference in jet axis position angle between the observations at these higher frequencies and the VSOP result at 5~GHz \\citep{Lis2001a}.  A careful inspection of our model-fitting results by allowing the change in PA shows that the uncertainties in the determination of the position angle for the inner 2 mas components are less than a few degrees. We also note that our model fitting results at 15 and 22~GHz are in good agreement with those at 8~GHz \\citep{Fey97}. All this indicates that the offset of the jet axis at 5~GHz and higher frequencies is real.  Frequency-dependent offsets in the position of jet axis at the VLBI scale have been seen in 3C~454.3 at frequencies 15 and 86~GHz \\citep{Kric99} and could be caused by the opacity effects enhanced in the regions of stronger jet bending. If the spectral index of the emission region in the transversal direction of the jet is different, the position of the fitted component may shift transversally with the observing frequency. Alternatively, a frequency-dependent core position offset may explain this difference \\citep{Loba98}. We also note that the restoring beam of the VSOP observation is 1.08~mas~$\\times$~0.22~mas at ${\\rm PA} = 20\\degr$, and the major axis of the beam is oriented along the offset direction of the jet axis. This alignment could introduce the apparent offset.  In addition to the 22 and 15~GHz VLBA data sets described above, we have got an access to the 5~GHz VLBA data from the VLBApls observations \\citep{Foma2000}. The image of J1625+4134 from VLBApls was originally model-fitted with two components \\citep{Foma2000}. In order to compare the models at 5~GHz and higher frequencies, we fitted the self-calibrated VLBApls $u,v$-data with a 5-component model which includes the northern component B1 detected in both images at 2.3 and 8.5~GHz \\citep{Fey97}. The new model-fitting results are given in Table~1. Fig.~3 shows the new model-fitting image based on the VLBIpls data which is consistent with the image published by Fomalont et al. \\citep{Foma2000}. The positions of the five model components for the VLBApls 5~GHz data are also shown in Fig.~2. Since the resolution at 5~GHz is lower than that at high frequencies and the structure of the original image at low resolution is simple, the reduced $\\chi^2$ of this model-fitting result (for 30 seconds averaged data) is relatively high, 2.29. The new model result cannot be considered as unique, it represents only one of many possible fits.  \\subsection {Variations of VLBI structure}  The C2 component is present in the VLBI models at 5 (both VSOP and VLBApls observations), 8, 15 and 22~GHz. It is located at the same position angle (about $-110\\degr$) in all observations except the VSOP one. In the latter, the component C2 is some $15\\degr$ off the position in all other data sets. If the C2 component in all the observations indeed corresponds to the same physical component, its apparent proper motion is $\\mu$ = 0.11$\\pm$0.04 mas/yr, corresponding to a $(7.4\\pm2.7\\,)c$ apparent transverse velocity. The position of component C2 as a function of time and the best fit of its apparent proper motion is shown in Fig.~4. The estimate of proper motion should be treated with caution due to possible opacity effects and spectral variations across the emission region.   \\subsection {Spectra of the components}  VLBI data on the source J1625+4134 mentioned above have been obtained at different epochs (2.3 and 8.5~GHz -- at 1995.77; 5~GHz -- 1996.42 and 1998.1; 15~GHz -- 1999.0; 22~GHz -- 2000.16). The UMRAO monitoring program at 4.8, 8.0 and 14.5~GHz shows that the total flux density of the source varied for up to 30\\% during the period 1995--2000. In spite of the source's variability and non-simultaneous VLBI data at different frequencies we attempted to compose and analyze the spectra of the components. The spectra are shown in Fig.~5. Although these spectra are no more than very rough estimates, they offer some useful insight into the radio emission from the jet components.  The component D has a flat spectrum with $\\alpha = 0.15\\pm0.26$ between 5 and 22~GHz. The higher flux density at 2.3~GHz may be due to the lower resolution which may allow for a contribution from low brightness extended emission surrounding the source. Based on its spectrum, the component D is the most likely candidate for the 'core' of the source.  The components C2, C2B and B1 have steep spectra. The spectral index of the component C between 5 and 22 GHz is $\\alpha=-1.3\\pm0.3$ (the VSOP data is not included in this spectral index estimate). The component C2B has $\\alpha= -1.2\\pm0.3$ between 8 and 22~GHz, its turnover frequency is probably between 5 and 8~GHz. The northern component B1 shows similar spectrum with the component C2B, $\\alpha = -1.1\\pm0.3$ between 2 and 15~GHz.  To obtain a two-frequency non-simultaneous spectral index map of the area around the core of J1625+4134, we produced the images at 15 and 22~GHz using the same $u,v$-ranges, cell size and restoring beams. The problem of producing a spectral index map is in determining the reference position in two images. We have decided to use the position of peak emission of the core component D to align the two images. An eastward shift of 0.04 mas (0.6 pixels) in right ascension was needed in the image at 22 GHz to align its D component with its counterpart at 15~GHz. We note that due to the possible self-absorption in the core region, the peaks of emission may lie in different parts of the core at different frequencies. However, the frequency difference between 15 and 22~GHz is relatively small. If the absolute core position is $r\\propto\\nu^{-1}$ \\citep{Loba98}, the estimated frequency-dependent position difference of the core in J1625+4134 is $< 0.1$~mas. Fig.~6 shows the spectral index distribution superimposed with the contours of the 15~GHz image.  As one would expect, in spite of an unknown impact of the source variability, the core area of the source has a flat spectrum. The jet region has mostly steeper spectrum, with two noticeable exceptions associated with the component C2 and the area between the components C2B and C1. The flatter spectrum around component C2 is likely to reflect the motion of the component outwards: the leading edge of the component is brighter at higher frequencies resulting in the flatter overall component spectrum. The second region of a flatter spectrum between the components C2B and C1 might be related to the structure variation or to the true spectral index distribution as well as the structural change as in the case of the component C2.   \\subsection {The brightness temperature of the mas-scale jet components}  The brightness temperature $T_{b}$ of an elliptic Gaussian component in the source rest frame is given by \\citep{Shen97}  \\begin{equation} T_{b} =1.22\\times 10^{12}(1+z)\\frac{{S_\\nu}}{{\\nu_{ob}^{2} ab}}, \\end{equation}  \\noindent where $S_\\nu$ is the flux density of the component in Jy at the observing frequency $\\nu$ in GHz, $a$ and $b$ are the major and minor axes in mas respectively, and $z$ is the source redshift. The brightness temperature distribution of the VLBI jet components, measured at the different frequencies, is shown in Fig.~7. Due to the resolution limit, the brightness temperature of the core may be only a lower limit and the source variation also may introduce some uncertainties in the estimate of the brightness temperatures.  The brightness temperature of the core is plotted in both panels of Fig.~7.  The high brightness temperature of the component D also supports its identification as the core. The brightness temperature decreases sharply within about 4~mas from the core in the southwest region. However, the brightness temperature in the northern jet is relatively constant from 4 to 25~mas. The decrease of the brightness temperature in the inner jet may be due to radiation losses and bending of the jet. If the brightness temperature variations are dominated by the variation of the viewing angle, the jet bending reduces the Doppler boosting effect in the southwest jet, while the orientation of the jet remains relatively unchanged in the northern region.   \\section {Discussion}  \\subsection {Spectral energy distribution of J1625+4134}  Using the data available in NED, we estimated the source rest-frame spectral energy distribution (SED). An upper limit flux density of 0.85~$\\mu$Jy at 1~keV was measured by HEAO-A1 \\citep{Bie87}. In the ROSAT-FIRST correlation study, Brinkmann et al.(2000) found 5 radio sources within the resolution of ROSAT around J1625+4134. Of these five, J1625+4134 is the nearest one to the ROSAT pointing center (the position offset is $5\\arcsec$) . The X-ray flux measured by ROSAT from this area is $0.14\\pm0.05 \\times 10^{-12} \\,{\\rm erg\\,s}^{-1}\\,{\\rm cm}^{-2}$. The information on the photon index for this source is unavailable, thus we have used the mean value of 2.46 for the quasars \\citep{Brink2000}. This leads to the estimated flux density of 0.012 $\\mu$Jy at 1~keV.  The peak frequency of the synchrotron emission energy, $\\nu_{\\rm peak} = 3.4\\times 10^{12}$~Hz, is located in the sub-mm band, the spectral energy at the peak frequency, $\\nu L_{\\nu} = 8.8\\times 10^{38}$~W, is typical for the radio-loud quasars \\citep{Sam96}. The X-ray emission could not fit the synchrotron emission, most likely this is to be attributed to the Self-Synchrotron Compton scattering (SSC).   \\subsection {Doppler boosting}  J1625+4134 is a typical radio loud quasar which shows all the indications on the relativistic bulk motion of the plasma in its jet, with its syncrotron emission Doppler boosted. There are several methods of estimating the Doppler boosting factor.  Assuming that the X-ray emission mechanism is SSC scattering, one can adopt the equation (1) of Ghisellini et al. \\citep{Ghis93} to estimate the Doppler boosting factor $\\delta_{sp}$ in a uniform sphere of the emission region. Readhead \\citep{Read94} suggested another method of estimating the Doppler boosting factor, the so called equipartition Doppler boosting factor $\\delta_{eq}$ . The method is based on the assumption on the energy equipartition between the radiating particles and the magnetic field.  Both methods above require VLBI core data at the turnover frequency. It proved to be difficult since the core has a very flat spectrum (see Fig.~5). In Table~2, we list the Doppler boosting factor estimates using both methods by assuming VLBI observing frequencies as the VLBI core turnover frequency. To avoid a fitted zero axial ratio for the 8.5 GHz core \\citep{Fey97}, we used a value of 0.2 for this ratio in the calculation.  In the equipartition case, a correction of the core size, $\\theta_{\\rm d}=1.8 \\sqrt{ab}$ was applied \\citep{Mar87, Gui96}, while in the $\\delta_{\\rm sp}$ estimate, a factor of 0.8 was adopted \\citep{Gui96}. The new X-ray flux density was used in the $\\delta_{\\rm sp}$ estimate and the spectral index $\\alpha = - 0.75$ for the optically thin synchrotron emission was assumed. Table~2 gives the estimates of the Doppler boosting factor for VLBI core data obtained at different frequencies.  Jiang et al. \\citep{Jiang98} investigated the inhomogeneous jet parameters in the K\\\"onigl's model \\citep{Koni81}. In this model, a VLBI core is believed to be a base of the optically thick synchrotron emission region of an inhomogenous jet. Assuming that the X-ray emission mechanism is the SSC scattering and the distributions of the magnetic field and the relativistic electron number density in the jet follow the power law ($m = 1$ for the magnetic field and $n = 2$ for the electron number density), and using the value of the proper motion of $\\mu = 0.11$~mas/yr, we can obtain the Doppler boosting factor and estimate other parameters of the model. The Doppler boosting factors $\\delta_{\\rm jet}$ for the inhomogeneous jet model are also given in Table~2.  These estimates of the Doppler boosting factor values cover a wide range of VLBI core parameters obtained at different frequencies. We assume the mean value $\\delta = 5.2\\pm 2.5$ as the best estimate of the Doppler boosting factor in the jet of J1625+4134. Combined with the proper motion mentioned above, the viewing angle to the line of sight could be estimated as $\\theta = 10\\degr \\pm 4\\degr$, and the Lorentz factor of the jet $\\gamma = 8 \\pm 4$. These parameters are consistent with those for radio-loud quasars.    \\subsection {Possible explanation for the jet structure}  The proper motion, non-detection of the counter-jet and the estimated value of the Doppler boosting factor suggest that the orientation of the inner south-west jet is close to the line of sight. These arguments rule out a two-sided jet morphology as an explanation of the observed bi-directional (southwest and north-northwest) jet structure in J1625+4134.  In radio-loud AGNs, large apparent bending angles of jet can be explained by the projection effect, since the viewing angle to the line of sight is small. The observed difference in the position angle of different areas along the jet, $\\Delta PA$, is related to the true bending angle of the jet $\\Delta \\phi$ as  \\begin{equation} cos{\\Delta\\phi}=cos\\theta_{1}cos\\theta_{2}+sin\\theta_{1}sin\\theta_{2}cos\\Delta {PA}, \\end{equation}  \\noindent where $\\theta_{1}$ and $\\theta_{2}$ are the viewing angles of the jet. The equation (2) holds only for the step change of the $\\Delta \\phi$. For a curved jet, the accumulated true bending angle $\\Delta \\phi$ should be integrated along the trajectory of the bending using the small angle approximation. If the jet moves along a cone with the half-opening angle $\\theta_{1}$ (this means that $\\theta_{2}$ is equal to $\\theta_{1}$), the integrated true bending angle $\\Delta \\phi = sin\\theta_{1}(\\Delta PA)$. In the J1625+4134 case, the estimated viewing angle of the southwest inner jet $\\theta_{1}$ is about $10\\degr$. If the two-sided jet is connected via an invisible curved trajectory, $\\Delta PA = 270 \\degr$ corresponds to $\\Delta \\phi = 47 \\degr$. This is a large intrinsic bending.  The cause for such a large bending of the jet in J1625+4134 is not clear. There is no obvious evidence of the enhanced emission at the region of sharp bending region of its jet. It seems that the bending of the jet is not caused by the interaction with the surrounding medium, since such interaction will produce strong shocks, further enhancing the radio emission.  Another possible explanation for the bi-directional jet structure is the intrinsic change in the jet ejection angle caused by a ballistic precession or erratic change in the nozzle direction. The bi-directional jet structure in J1625+4134 is more like the erratic change in the nozzle direction, because two jets follow two very different ballistic trajectories \\citep{Gome99}. The northern jet may have been ejected when the nozzle of the jet pointed to the north and the southwest jet is the result of a newer ejection. In this case, the change of the nozzle direction would occur more than about 50 years ago -- the time needed for the bending southwest jet material to reach its present extension of about 4-5~mas with the apparent proper motion of the component C2 of about 0.1~mas/yr.  Since the northern emission has a steep spectrum, a high resolution observations with a higher sensitivity at low frequencies may be useful to distinguish between the two possible models. A detection of the connection between the two jets will support the projection effect explanation. ",
        "conclusions": "The quasar J1625+4134 shows an uncommon bi-directional jet structure on the VLBI scale at various observing frequencies. We presented the images of VLBA observations at 15 and 22~GHz, showing a short curved southwest jet which extends for several mas from the core. Our VLBI results, combined with the data available in the literature, allow us to estimate the spectral index of the components and to identify the component D as the VLBI core based on its flat spectrum between 5 and 22~GHz. The jet components distribution at different observing frequencies favors a large difference of about $270 \\degr$ in the position angles of the direction of the jet propagation in the southwest and northern areas from the core. A proper motion of the component C2 in the southwest jet, $\\mu = 0.11$~mas/yr, was estimated. Using the multi-band VLBI data on the component D and different physical methods, we obtained the best estimate of the Doppler boosting factor, $\\delta =  5.2\\pm 2.5$, corresponding to the southwest inner jet viewing angle of $10\\degr\\pm 4\\degr$ to the line of sight with the Lorentz factor $\\gamma = 8\\pm 4$.  The proper motion, the estimated Doppler boosting factor and non-detection of the counter-jet emission at the higher VLBI frequencies rule out the two-sided-jet explanation for the observed bi-directional jet morphology. According to the analysis on the relativistic jet parameter, we point out on the possibility that the southwest jet goes to the north via a helical trajectory, while an erratic change in the nozzle direction also can explain the bi-directional jet structure. High resolution VLBI observations at lower frequencies may provide clue on this model.  We found a difference of about $15\\degr$ in the southwest jet axis between the 5~GHz VSOP image and other VLBA observations at higher frequencies. Our re-analysis of the 5~GHz VLBApls data narrows the difference in the axis position, but the hypothesis on a frequency dependent position of the jet axis needs further verification."
    },
    "0206/astro-ph0206365_arXiv.txt": {
        "abstract": "We present results from the first hydrodynamical star formation calculation to demonstrate that brown dwarfs are a natural and frequent product of the collapse and fragmentation of a turbulent molecular cloud.  The brown dwarfs form via the fragmentation of dense molecular gas in unstable multiple systems and are ejected from the dense gas before they have been able to accrete to stellar masses. Thus, they can be viewed as `failed stars'.  Approximately three quarters of the brown dwarfs form in gravitationally-unstable circumstellar discs while the remainder form in collapsing filaments of molecular gas.  These formation mechanisms are very efficient, producing roughly the same number of brown dwarfs as stars, in agreement with recent observations. However, because close dynamical interactions are involved in their formation, we find a very low frequency of binary brown dwarf systems ($\\lsim 5$\\%) and that those binary brown dwarf systems that do exist must be close $\\lsim 10$ AU.  Similarly, we find that young brown dwarfs with large circumstellar discs (radii $\\gsim 10$ AU) are rare ($\\approx 5$\\%). ",
        "introduction": "The existence of brown dwarfs was incontrovertibly demonstrated for the first time by the discovery of Gliese 229B \\cite{Nakajimaetal1995}, a cool brown dwarf orbiting an M-dwarf. In the same year, other candidates later confirmed to be free-floating brown dwarfs were announced (e.g.\\ Teide 1 by Rebolo, Zapatero-Osorio \\& Martin 1995), along with PPl 15 which was later discovered to be a binary brown dwarf (Basri \\& Martin 1999). Observations now suggest that brown dwarfs are as common as stars, although stars dominate in terms of mass (e.g.\\ Reid et al.\\ 1999).  Despite the abundance of brown dwarfs, their formation mechanism is currently a mystery.  The typical thermal Jeans mass in molecular cloud cores is $\\approx 1$ M$_\\odot$ (Larson 1999 and references therein). Thus, the gravitational collapse of these cores might be expected to form stars, not brown dwarfs.  There are two obvious routes by which brown dwarf systems (i.e.\\ brown dwarfs without stellar companions) may form. First, they may result from the collapse of low-mass cores (masses $\\lsim 0.1$ M$_\\odot$) that are smaller (radii $\\lsim 0.05$ pc) and denser ($n$(H$_2$)$\\gsim 10^7$ cm$^{-3}$) than the cores that are typically observed (i.e.\\ they have low masses yet are still Jeans unstable). Thus, brown dwarfs would be `low-mass stars'.  Such low-mass bound cores have not yet been observed, but they would be difficult to detect because of their small sizes and low masses, and rarely detected due to their short dynamical time-scales ($\\sim 10^4$ years). Observations are beginning to reach this mass regime (e.g.\\ Motte, Andr\\'e \\& Neri 1998), although the low-mass clumps found thus far probably are not gravitationally bound \\cite{Johnstoneetal2000}.  The second possibility is that brown dwarfs form in higher-mass cores but are prevented from accreting enough mass to exceed the hydrogen-burning limit.  If such a core fragments to form an unstable multiple system, this may be achieved by the dynamical ejection of a fragment from the core, cutting it off from the reservoir of gas, and thus preventing it from accreting to a stellar mass. In this case, brown dwarfs would be `failed stars'. This ejection mechanism has been proposed by Reipurth \\& Clarke \\shortcite{ReiCla2001} and Watkins et al.\\ \\shortcite{Watkinsetal1998b}.  In this letter, we present results from the first hydrodynamical calculation to demonstrate that a large number of brown dwarfs can be formed during the fragmentation of a molecular cloud. All the brown dwarfs are formed by the ejection of fragments from unstable multiple systems. In section 2, we briefly describe the numerical method and the initial conditions for our calculation.  In section 3, we present results from our calculation and compare them with observations.  Finally, in section 4, we give our conclusions. ",
        "conclusions": "We have presented results from the first calculation to demonstrate the formation of brown dwarfs from the fragmentation of a turbulent molecular cloud.  The calculation resolves the opacity limit for fragmentation and follows numerous brown dwarfs until accretion onto them has ceased.  We find that the star-formation process produces roughly equal numbers of stars and brown dwarfs, in agreement with recent observations.  Examining the mechanisms by which the brown dwarfs form, we find that they are the result of the ejection of fragments from the dense gas in which they form by dynamical interactions in unstable multiple systems.  This occurs before they are able to accrete to stellar masses. Three quarters of the brown dwarfs fragment out of gravitationally-unstable discs, while the remainder form in collapsing filamentary flows of high-density gas.  The calculation indicates that close binary brown dwarf systems (separations $\\lsim 10$ AU) might be able to survive the ejection process.  However, such systems should be very rare (frequency $\\lsim 5$\\%) because of the close dynamical interactions that are involved in the ejection of brown dwarfs.  Similarly, large discs (radii $\\gsim 10$ AU) around young brown dwarfs should also be rare (frequency $\\approx 5$\\%).  Observations show that close binary brown dwarfs and circumstellar discs surrounding brown dwarfs do exist, but more detailed surveys are required to test the predictions made here."
    },
    "0206/astro-ph0206153_arXiv.txt": {
        "abstract": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}  Spectroscopy with the Keck~II 10-meter telescope\\footnote{Data presented herein were obtained at the W.\\ M.\\ Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration.  The Observatory was made possible by the generous financial support of the W.\\ M.\\ Keck Foundation.} and Echelle Spectrograph and Imager is presented for six Virgo Cluster dwarf elliptical (dE) galaxies in the absolute magnitude range $-15.7\\le{M_V}\\le-17.2$.  The mean line-of-sight velocity and velocity dispersion are resolved as a function of radius along the major axis of each galaxy, nearly doubling the total number of dEs with spatially-resolved stellar kinematics.  None of the observed objects shows evidence of strong rotation: upper limits on $v_{\\rm rot}/\\sigma$, the ratio of the maximum rotational velocity to the mean velocity dispersion, are well below those expected for rotationally-flattened objects.  Such limits place strong constraints on dE galaxy formation models.  Although these galaxies continue the trend of low rotation velocities observed in Local Group dEs, they are in contrast to recent observations of large rotation velocities in slightly brighter cluster dEs.  Using surface photometry from {\\it Hubble Space Telescope\\/}\\footnote{Based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS~5-26555.} Wide Field Planetary Camera~2 images and spherically-symmetric dynamical models, we determine global mass-to-light ratios $3\\le\\Upsilon_V\\le6$.  These ratios are comparable to those expected for an old to intermediate-age stellar population and are broadly consistent with the observed $(V-I)$ colors of the galaxies.  These dE galaxies therefore do not require a significant dark matter component inside an effective radius.  We are able to rule out central black holes more massive than $\\sim10^7\\Msun$.  For the five nucleated dEs in our sample, kinematic and photometric properties were determined for the central nucleus separately from the underlying host dE galaxy.  These nuclei are as bright or brighter than the most luminous Galactic globular clusters and lie near the region of Fundamental Plane space occupied by globular clusters.  In this space, the Virgo dE galaxies lie in the same general region as Local Group and other nearby dEs, although non-rotating dEs appear to have a slightly higher mean mass and mass-to-light ratio than rotating dEs; the dE galaxies occupy a plane parallel to, but offset from, that occupied by normal elliptical galaxies. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}  Dwarf elliptical galaxies (dEs) are the most common galaxy type by number in the Local Universe, dominating the galaxy luminosity function of nearby clusters.  Yet these galaxies remain among the most poorly studied galaxies due to their faint luminosities, $M_V \\ge -18$, and characteristic low effective surface brightness $\\mu_{V,\\rm eff}>22$~mag~arcsec$^{-2}$ \\citep{fer94}.  Unlike brighter, classical elliptical galaxies whose surface brightness profiles are well fit by the de~Vaucouleurs $r^{1/4}$ law \\citep{dev48}, dEs have brightness profiles that are characterized by Sersic profiles \\citep{ser68} with indices ranging between $n=1$--3 (where $n=1$ corresponds to an exponential law and $n=4$ to an $r^{1/4}$ law) making them appear more diffuse than classical ellipticals of the same total magnitude \\citep{bin98}.  In the Virgo Cluster, the majority of dEs brighter than $M_V \\simlt -16$ contain compact central nuclei; fainter than $M_V \\simgt -12$ most dEs show no sign of a nucleus \\citep{san85}. Nuclei typically contain 5\\% to 20\\% of the total galaxy light and are slightly resolved at the distance of Virgo by {\\it Hubble Space Telescope\\/} ({\\it HST\\/}) imaging \\citep{mil98}.  In hierarchical models of galaxy formation, dwarf galaxies form out of small density fluctuations in the early Universe and are predicted to be less spatially clustered than normal elliptical or spiral galaxies \\citep{dek86}.  However, dwarf elliptical galaxies are preferentially found in dense cluster environments, more so than either ellipticals or spirals \\citep{bts87}; there are few, if any, examples of isolated dEs.  Thus, current models favor dE formation from a progenitor galaxy population.  The proposed progenitors of dEs are spiral or irregular galaxies which are morphologically transformed into dEs through the processes of galaxy harassment and interaction \\citep*{moo98}. Detailed internal kinematics of dEs are a powerful observational tool with which to test these scenarios.  Until recently, radial velocity dispersion profiles were only available for the Local Group dEs and two of the brightest dEs in the Virgo Cluster \\citep*{ben90,ben91,hel90}.  In addition, a handful of global velocity dispersion measurements existed for dEs in various environments outside the Local Group \\citep{pet93}.  These observations suggest that dEs have lower mass-to-light ratios than Local Group dwarf spheroidals (e.g.,~Draco, Fornax) and are flattened by velocity anisotropy rather than by rotation.  However, \\citet{ped02} and \\citet{der01} recently presented kinematic profiles for a few dEs in the Virgo and Fornax Clusters, respectively, with rotation velocities comparable to that expected for a rotationally-flattened spheroid.  These rotating dEs are more luminous on average than the non-rotating dEs we have observed, and hint at a possible association between dE luminosity and the presence of rotation.  The question of whether dEs have significant rotation compared to their velocity dispersion is particularly important in the context of dE formation scenarios.  \\citet{moo98} have demonstrated that the process of galaxy harassment in cluster environments can morphologically transform a spiral galaxy into a dE.  Although this process tends to increase the velocity dispersion in a system, it is less efficient at disrupting rotational motions and a significant fraction of the progenitor's rotation is preserved.  Thus, measuring the amount of internal angular momentum in dEs can constrain the progenitor galaxy type and/or the amount of disruption required.  We present internal kinematics as a function of radius for a sample of six~dE galaxies in the Virgo Cluster based on Keck observations. These data are interpreted in conjunction with archival {\\it HST\\/} imaging.  Preliminary results from this study were presented in \\citet*{geh02}.  This paper is organized as follows: in \\S\\,\\ref{data_sec} we present the Keck spectroscopic and {\\it HST\\/} imaging observations of our target galaxies, along with an outline of data reduction procedures; in \\S\\,\\ref{res_sec} we present velocity and velocity dispersion profiles and describe the dynamical models that are applied to the data to derive mass-to-light ratios, constraints on orbital anisotropy, and limits on the central black hole mass; the broader implications of the results are discussed in \\S\\,\\ref{disc_sec}. ",
        "conclusions": "\\label{disc_sec}  Velocity and velocity dispersion profiles are presented for the major axes of six dE galaxies in the Virgo Cluster.  These galaxies do not show evidence for substantial rotation; upper limits on rotation velocities are well below that expected if these objects were rotationally flattened.  Dynamical models for these galaxies suggest mass-to-light ratios in the range $3\\le \\Upsilon_V \\le6$.  We argue that such ratios are expected for intermediate to old stellar populations and thus these dEs do not require significant dark matter inside an effective radius.  Our observations do not rule out significant dark matter in dEs at larger radii as demonstrated by giant elliptical galaxies which exist in massive dark halos, but are not necessarily dark matter dominated at small radii \\citep{ger01}. In Fundamental Plane space, we find that the Virgo dE galaxies, similar to previously observed dEs, lie in a plane parallel to, but offset from, that occupied by normal elliptical galaxies.  In this space, dE nuclei lie near the region occupied by Galactic globular clusters.  The origin of nuclei in dE galaxies remains an open question.  In the present sample, there is no obvious difference between the single non-nucleated dE galaxy (VCC~917) and the underlying galaxies of the observed nucleated dEs.  The mass-to-light ratio, anisotropy, and photometric parameters measured for VCC~917 are indistinguishable from those determined outside the nucleus of the other five dEs.  However, as a population, non-nucleated dE galaxies in the Virgo Cluster do have different properties.  They are observed to be less spatially concentrated, have lower specific globular cluster frequencies, and, on average, have flatter shapes as compared to nucleated dE galaxies \\citep{san85,mil98,ryd99}.  Proposed scenarios for the origin of dE nuclei include the remnant cores of larger stripped galaxies \\citep{ger83}, the results of gas infall and star formation or the coalescence of several globular clusters whose orbits have decayed to the dE center \\citep{oh00}.  We have shown that the observed dE nuclei share many properties with globular clusters, suggesting similar formation processes.  Since dE galaxies are preferentially found in dense environments, it is likely that galaxy interactions play a large role in their formation and evolution.  The models of \\citet{moo98} suggest that galaxy harassment in clusters can morphologically transform a spiral galaxy into a dwarf elliptical.  Harassment tends to increase internal velocity dispersions, but is less efficient in disrupting rotational motion and is not obviously reconciled with the low rotational velocities observed in the present dE sample.  If dEs are the morphologically-transformed remnants of larger progenitor galaxies, a constraint on such a progenitor population is provided by the central black hole mass limits determined in \\S\\,\\ref{bh_sec}.  The upper limit of $\\sim10^7\\Msun$ for the observed dE galaxies implies that any dE progenitor must have had a bulge dispersion less then 100\\kms, assuming the $M_{\\rm BH}-\\sigma_e$ relation \\citep{tre02}.  Although this is not a stringent constraint on dE galaxy formation models, higher spatial resolution kinematics, and therefore more stringent mass limits, could be a significant constraint on such models.  As the number of dE galaxies with measured internal kinematics increases, their position in the Fundamental Plane strengthens the conclusion that dwarf and classical elliptical galaxies evolve via very different physical processes.  To determine whether dwarf ellipticals as a galaxy class evolve under homogeneous conditions requires more observations.  A critical question is understanding the apparent dichotomy between the anisotropy-supported dEs presented in this paper and the rotationally-supported dEs presented by \\citet{ped02} and \\citet{der01}.  The fact that rotating and non-rotating dEs appear to form a ``sequence'' in Fundamental Plane space, with the latter having somewhat lower mean luminosity, mass, and mass-to-light ratio, suggests that these are not two distinct types of dE galaxies but rather are part of a continuous family. Larger samples are required to establish what, if any, physical property correlates with the observed rotational velocities and what this implies for dE galaxy formation."
    },
    "0206/astro-ph0206479_arXiv.txt": {
        "abstract": "We review the behavior of the oscillating shear layer produced by gravity waves below the surface convection zone of the Sun. We show that, under asymmetric filtering produced by this layer, gravity waves of low spherical order, which are stochastically excited at the base of the convection zone of late type stars, can extract angular momentum from their radiative interior. The time-scale for this momentum extraction in a Sun-like star is of the order of $10^7$ years. The process is particularly efficient in the central region, and it could produce there a slowly rotating core. ",
        "introduction": "Angular momentum transport by gravity waves has received much attention recently. While it has first been considered as a key mechanism in the tidal interaction of binary systems (Zahn 1975; Goldreich \\& Nicholson 1989), it has since been invoked also as an efficient process of momentum redistribution in single stars (Schatzman 1993; Kumar \\& Quataert 1997; Zahn, Talon \\& Matias 1997). The mechanism proposed was similar to that acting in binary stars, with synchronization  proceeding gradually inwards. However, as pointed out by Gough \\& McIntyre (1998) and Ringot (1998), the treatment of angular momentum extraction as presented by these first studies was incorrect, and it became clear that the actual properties of wave transport were far more complex than originally anticipated.  Wave properties have been examined further by Barnes, McGregor \\& Charbonneau (1998) who showed that magnetic fields transform pure gravity waves into gravito-Alfv\\'en waves, thus modifying their damping behavior. They showed that strong magnetic field may prevent certain waves from propagating.  Kumar, Talon \\& Zahn (1999, hereafter KTZ) then demonstrated how gravity waves in interaction with shear turbulence may lead to the formation of an oscillating shear layer just below the surface convection zone, analogous to the well studied quasi-biennial oscillation of atmospheric sciences (see \\eg Shepherd 2000 for details on the QBO and other wave-driven oscillations).  Kim \\& McGregor (2001) further studied that oscillation in a simplified two waves model. They showed that, with only a prograde and a retrograde wave, and with the velocity fixed both at the top and at the bottom of the solar tachocline, the time-dependent behavior of the rotation profile goes from being periodic, to quasi-periodic and to chaotic as the viscosity in the shear region is slowly decreased. Their goal was to seek an explanation of a 1.3 yr variation, which Howe et al. (2000) claim to have detected in the tachocline.  Here we want to reexamine the effect of waves in the deep solar interior. Indeed, while the high degree waves are damped very close to the base of the convection zone, thus leading to rapid oscillations that could explain rotational velocity variations in the tachocline (Kim \\& McGregor 2001) or be related to the solar cycle (KTZ), the low degree waves that are also excited stochastically by the convective motions, although with a somewhat lower efficiency (see Kumar \\& Quataert 1997 and Zahn, Talon \\& Matias 1997) will have, over a much longer period, an effect on the deep interior.  In this letter, we wish to demonstrate how these waves, in conjunction with shear turbulence, indeed achieve momentum extraction in the deep interior. In \\S \\ref{shme}, we explain heuristically how that process takes places and in \\S \\ref{nums}, we show the results of a long numerical simulation applied to the solar case.  \\begin{figure}[t] {\\centering \\resizebox*{0.46 \\textwidth}{!} {\\rotatebox{0}{\\includegraphics{spectre.ps}}} \\par} \\caption{Angular momentum luminosity integrated over $0.1 \\mu{\\rm Hz}$ as a function of order $\\ell$ and frequency. \\label{fig:mom_lum}} \\end{figure}   \\begin{figure}[t] {\\centering \\resizebox*{0.46 \\textwidth}{!} {\\rotatebox{-90}{\\includegraphics{local3.ps}}} \\par} \\caption{Oscillating shear layer below the surface convection zone. The dotted line shows the initial rotation profile. With the surface rotating slower than the core, a prograde shear layer is initially formed, followed by a retrograde one. When the shear becomes too intense, turbulent viscosity acts to merge the prograde layer with the convection zone, leaving behind the retrograde layer. A new prograde layer forms behind, and migrates towards the convection zone when the retrograde layer is  absorbed. The cycle then resumes. \\label{fig:local}} \\end{figure}   \\begin{figure}[t] {\\centering \\resizebox*{0.46 \\textwidth}{!} {\\rotatebox{0} {\\includegraphics{global.ps}}} \\par} \\caption{Global evolution of the rotation profile over large time periods. Efficient momentum extraction is visible in a characteristic time-scale of $10^5$ years. \\label{fig:global}} \\end{figure}   \\begin{table}[t] \\caption{Damping length $d_\\omega$  (scaled by the solar radius) at $0.6 ~ \\mu {\\rm Hz}$} \\begin{minipage}{\\linewidth} \\begin{tabular}{rrrr}\\hline \\multicolumn{1}{c}{$\\ell$} & \\multicolumn{1}{c}{$m$} & \\multicolumn{1}{c}{$d_\\omega$\\footnote{In solid body rotation}} & \\multicolumn{1}{c}{$d_\\omega$\\footnote{With the initial rotation profile (\\cf Figures~\\ref{fig:local} \\&~\\ref{fig:global})}} \\\\ \\hline \\hline 1  &  1  & 0.616 & 0.461 \\\\ & -1  &       & 0.661\\vspace{0.1cm}\\\\ 5  &  5  & 0.084 & 0.036 \\\\ & -5  &       & 0.499\\vspace{0.1cm}\\\\ 9  &  9  & 0.033 & 0.017 \\\\ & -9  &       & 0.378\\vspace{0.1cm}\\\\ 13  &  13 & 0.018 & 0.010 \\\\ & -13 &       & 0.256\\vspace{0.1cm}\\\\ 17  &  17 & 0.012 & 0.007 \\\\ & -17 &       & 0.061\\vspace{0.1cm}\\\\ 21  &  21 & 0.008 & 0.005 \\\\ & -21 &       & 0.024\\vspace{0.1cm}\\\\ \\end{tabular} \\end{minipage}\\label{tab:damp} \\end{table} ",
        "conclusions": "In this letter, we have shown how gravity waves can extract angular momentum from the radiative core of a solar-like star through differential wave filtering, sufficiently to establish an almost uniform rotation profile in about $10^7$ years. In our model, the presence of turbulent transport is essential to yield the proper result; indeed, it is required both to produce the QBO like oscillation, and to smooth the excessive differential rotation produced by waves in the central region.  The evolution of latitudinally averaged angular momentum in the Sun would then proceed as follows. The magnetic torque applied to the convection zone tends to slow it down, establishing a negative gradient of angular velocity, which induces the bias between prograde and retrograde waves. Due to the combined action of wave transport and of turbulent viscosity, the velocity gradient tends to diminish. In our simulation it disappears altogether, but in the real Sun, it would be maintained at some level due to the competition between the extraction of angular momentum from the radiative interior, and its loss by the solar wind: the faster the spin-down, the steeper the gradient, the stronger the extraction. If the extraction of angular momentum by gravity waves is as efficient as we estimate, it will adjust itself such as to provide just the flux of angular momentum which is lost by the wind, with the interior profile of angular velocity being rather flat, except near the center and at the top of the radiation zone.  The presence of a large toroidal magnetic at the base of the convection zone could somewhat lengthen the time-scales mentioned here. Indeed, the magnetic field required to prevent wave propagation at $0.6 ~{\\rm \\mu Hz}$ is about $(3 \\times 10^5/\\ell)~{\\rm G}$ (\\cf KTZ). If fields of a strength as high as $10^5 ~{\\rm G}$ are present in that region (\\cf Fan et al. 1993, Caligari et al. 1995), high $\\ell$ modes could be prevented from propagating, modifying the evolution of the double shear layer. However, since such large fields occur in localized areas, the global effect would rather be one of decreasing the wave flux and thus, simply increase the time scale estimated here.   This initial work must be pursued in many ways. Firstly, one would like to assess the efficiency of this mechanism in an evolving solar model, with magnetic spin-down, in which meridional circulation would also be taken into account in order to verify the scenario we proposed earlier.  Other improvements should be made on the physical description. The effect of the Coriolis force must be included, which could require to step up to two-dimensional simulations. Finally, a more realistic prescription for the turbulent transport must be implemented, before we can conclude that the solar core is rotating significantly slower than the rest of the radiation zone."
    },
    "0206/hep-ph0206185_arXiv.txt": {
        "abstract": "The process of the vacuum polarization energy losses of high energy cosmic rays propagating in the extragalactic space is considered. The process is  due to the polarization of Cosmic Background Radiation by a moving charged particle. With the goal of the description of the process, the photon mass, refractive indices and permittivity function for low and high energy photons are found. Calculations show the rather noticeable level of the energy losses for propagating protons with the energies more than $10^6 - 10^7 GeV$. The influence of the polarization energy losses on  propagation of cosmic rays is discussed. ",
        "introduction": "The problems of the origin and propagation of  high energy cosmic rays is widely discussed in recent years (see, for example \\cite{NW,BGG,AP} and literature therein). The photoproduction of mesons and $e^\\pm$-pairs in the Cosmic Background Radiation was considered as the main reason of the energy losses of charged particles propagating in the extragalactic space\\cite{BGG,Lee,St}. However, the simulations of the spectral distributions of cosmic rays on this basis do not provide a close  agreement with observed data \\cite{BGG}, and some peculiarities in the spectrum like the \"knee\" \\cite{K} and  GZK-cutoff \\cite{KG,ZK} do not find recognized explanation. In this connection the determination  of  another reason of the energy losses is important and interesting.   In this paper we consider the possible mechanism by which  cosmic rays lose their energy in the extragalactic space. It is the polarization energy losses  of a charged particle moving in the electromagnetic vacuum. In the presence of an external  electromagnetic field the polarization of vacuum was considered first in the pioneer papers \\cite{JS}. The vacuum polarization  leads  to different effects \\cite{LL1} such as the nonlinearity of the Maxwell equations, appearance of nonzero photon mass, birefringence of light, etc. The description of the various methods in the QED of vacuum one can find in literature ( see \\cite{FGS} and references therein). In parallel with other methods  the traditional method, based on the introduction of the permittivity tensor, may be used  for the effective description  of different phenomena in the electromagnetic vacuum. In this case the vacuum described by the Maxwell equations which are similar to equations of the classical electrodynamics for continuous media\\cite{LL,AG}. As an example, one can point out the well known  low energy permittivity and permeability tensors for constant electric and magnetic fields \\cite{LL1}. In paper \\cite{BK} the connection between the polarization  and permittivity tensors is established. The using of the permittivity tensor  is convenient because this approach allows to consider the QED vacuum and matter with the unit point of view. ",
        "conclusions": "On the  basic of determination  such characteristics as the photon mass, refractive indices and permittivity function in the Cosmic Background Radiation the vacuum polarization losses of high energy cosmic rays are considered. The calculations show the high level of these losses for protons with energies more than $\\approx 10^7 GeV$. The proposed mechanism of losses leads to a revision of the existing propagation models of cosmic rays. With our point of view the propagation of the high energy cosmic rays in the extragalactic space is the dynamic process to a greater extent than it is expected. Experimental and theoretical ivestigations of these processes will help to understand the nature and origin of the cosmic rays in the Universe.  The author would like to thank H. Zaraket for critical questions, remarks and useful  references."
    },
    "0206/astro-ph0206403_arXiv.txt": {
        "abstract": "We investigate the color-magnitude distribution in the rich cluster AC\\,118 at $z=0.31$. The sample is selected by the photometric redshift technique, allowing to study a wide range of properties of stellar populations, and is complete in the K-band, allowing to study these properties up to a given galaxy mass. We use galaxy templates based on population synthesis models to translate the physical properties of the stellar populations - formation epoch, time-scale of star formation, and metallicity - into observed magnitudes and colors.  The distributions of galaxies in color-magnitude space thus map into distributions in the space of physical parameters. This is achieved by means of a statistical procedure which constrains the photometric properties of AC\\,118 galaxies to reproduce those of a nearby rich cluster once evolved at $z\\sim0$. In this way we show that a sharp luminosity-metallicity relation is inferred without any assumption on the galaxy formation scenario (either monolithic or hierarchical).  Our data exclude significant differences in star formation histories along the color-magnitude relation, and therefore confirm a pure metallicity interpretation for its origin, with an early ($z\\sim5$) formation epoch for the bulk of stellar populations. The dispersion in the color-magnitude diagram implies that fainter galaxies in our sample (K$\\sim18$) ceased to form stars as late as $z\\sim 0.5$, in agreement with the picture that these galaxies were recently accreted into the cluster environment. The trend with redshift of the total stellar mass shows that half of the luminous mass in AC\\,118 was already formed at $z \\sim 2$, but also that 20\\% of the stars formed at $z<1$. ",
        "introduction": "The color-magnitude relation (CMR) of cluster early-type galaxies has been extensively investigated at $z < 1$ to trace their star formation history and hence to constrain their formation epoch (e.g. e.g Kodama \\& Arimoto 1997, hereafter KA97; Ellis et al. 1997; Gladders et al. 1998; Stanford, Eisenhardt, \\& Dickinson 1998; Kodama \\& Bower 2001, hereafter KB01; Smail et al. 2001; van Dokkum et al. 2001).  The most important observational results are: i) the slope of the CMR does not depend on redshift; ii) the optical-NIR rest-frame colors of early-type cluster members become bluer with increasing redshift; iii) the intrinsic scatter in the optical-NIR colors of early-type galaxies is small at all redshifts (e.g. Ellis et al. 1997; Stanford et al. 1998; van Dokkum et al. 1998, 2000; Kodama et al. 2001). These points lead to explain the color-magnitude (CM) sequence as a correlation between galaxy mass and metallicity, while the age of galaxies play only a marginal role, if any (e.g. KA97).  Two different scenarios can successfully explain the CMR as function of redshift: the monolithic collapse (e.g. Eggen, Lynden-Bell, \\& Sandage 1962; Tinsley \\& Gunn 1976) and the hierarchical merging (e.g. Kauffmann 1996, Kauffmann \\& Charlot 1998).  In the former, the trend of the mass-metallicity sequence is explained by the fact that the more massive galaxies retain supernova ejecta more effectively, resulting in higher metallicities and hence in redder colors for more luminous galaxies (e.g. Arimoto \\& Yoshii 1987; KA97). The unchanged scatter of the colors of early-type galaxies with redshift indicates either that the galaxies assembled synchronously over redshifts (at least for $z<1$) or that they stochastically formed at much earlier times (see Ellis et al. 1997).  For what concerns the alternative picture, Kauffmann \\& Charlot (1998) claimed that the CMR can be reproduced in a hierarchical merging picture, where the more massive/metal-rich ellipticals result from mergers of massive/metal-rich progenitor disk galaxies. In both scenarios the color evolution of early-type cluster galaxies is in agreement with the passive evolution of an old stellar population formed early in the past (see also Stanford et al. 1998; Kodama et al. 1998).  Both the main evolutionary scenarios have to face with the evidence for the presence of a significant population of blue galaxies in rich cluster environments at $z\\geq$ 0.2, as shown for the first time by Butcher \\& Oelmer (1978) and has confirmed by several photometric and spectroscopic observations (e.g. Butcher \\& Oelmer 1984; Couch \\& Newell 1984; Ellis et al. 1985; Dressler \\& Gunn 1982; Couch \\& Sharples 1987, Couch et al. 1994; Dressler et al. 1994).  Taking into account a representative sample of the whole cluster population, KB01 re-investigated the photometric Butcher-Oelmer (B-O) effect in distant clusters. They found that the passive evolution of galaxy populations can reconcile the B-O effect with the tight CMR of the Coma cluster. Furthermore, KB01 found that the distribution in the color-magnitude diagrams suggests a scenario where star formation of galaxies accreted by the cluster declines on a 1 Gyr time-scale and it is not sharply truncated by interaction with the cluster environment. In this scenario, the B-O effect depends on the decline of star formation of field galaxies when they are accreted into the cluster and on the decline of the rate of accretion of new galaxies at lower redshifts.  In the present work we will apply the CM diagram to gain insight into the star formation history in the galaxy cluster AC\\,118 at $z=0.31$. We will use population synthesis models in order to describe the observed CM distribution of galaxies in AC\\,118 in terms of stellar populations parameters. The cluster sample is selected according to the photometric redshift technique, and is complete in K-band, avoiding biases introduced by measuring the blue wavelengths in the cluster rest-frame.  The early-type galaxy population in the core of AC\\,118 was already analyzed by Stanford et al. (1998) who found evidence in favor of the passive evolution scenario.  A spectroscopic study of the cluster was performed by Couch \\& Sharples (1987) and Barger et al. (1996) who claimed for recent ($\\lesssim 2$ Gyr) bursts of star formation.  The layout of the paper is the following.  In \\S~2 we describe the sample of galaxies at $z\\sim 0.3$.  In \\S~3 we introduce the galaxy templates that will be used to interpret the observed photometry in terms of physical properties of stellar populations and we describe our approach. The resulting distribution of the physical parameters is analyzed in \\S~4, where we also discuss the origins of the CMR and the global star formation history.  In \\S~5 we summarize the main aspects of the work and draw the conclusions.  In the following we assume $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ and $\\mathrm{ H_0= 70~Km s^{-1} Mpc^{-1} }$. With this cosmology the age of the universe is $\\mathrm{13.5~Gyr}$, and the redshift of AC\\,118 corresponds to a look-back time of $\\mathrm{3.5~Gyr}$. We verified that changing the cosmology does not affect the results of the present work. ",
        "conclusions": "\\label{FINE}  We have studied the star formation history of galaxies in the rich cluster AC\\,118 at redshift $z=0.31$ by constraining their photometric properties to reproduce, once evolved at $z\\sim0$, those of a local rich cluster. The analysis is based on a large wavelength baseline including accurate VRIK photometry for a large sample of cluster galaxies (N=252). The sample was selected by the photometric redshift technique and is complete in the NIR, thus reducing possible biases towards objects with more recent/intense star formation activity.  One of the main current issues in the comparison of the properties of local and in\\-ter\\-me\\-dia\\-te--redshift clusters concerns the selection criteria of the samples.  Studies of the CM relation based on pure morphological selection can be biased towards the older progenitors of nearby early-type galaxies (see van Dokkum et al. 2000 for a detailed discussion).  On the other side, the application of a statistical field subtraction approach requires a wide area around the cluster field to be observed, while the use of a spectroscopically selected sample at faint luminosities is very expensive in terms of observing time (but not impossible, see Abraham et al. 1996; van Dokkum et al. 2000).  The main advantage of a selection based on photometric redshifts is that it allows to estimate the typical luminosity-weighted formation epoch of a stellar population irrespective of the past history of the host galaxy (such as, for example, clustering through a merging hierarchy), and it is therefore an ideal tool to define cluster membership for large samples of galaxies without any tie to a particular scenario of galaxy formation.  A more tricky point is represented by the areas of the clusters to be compared at different redshifts. In a hierarchical clustering picture, clusters of galaxies are likely to accrete a significant fraction of their population from the field even at relatively modest redshifts ($z<0.5$, see Kauffmann 1996). As a consequence, cluster richness tends to increase with time, while the population accreted at an old epoch becomes concentrated in a progressively smaller area (see e.g. KB01).  On the other hand, the cores of rich nearby clusters are very similar in their photometric properties: the galaxy population is dominated by E/S0 galaxies with few Sp/Irr having bluer colors. Moreover, early-type galaxies seem to follow a universal well defined CM relation (see Bower et al. 1992a, Bower et al. 1992b).  For such reasons, we have analyzed the constraints on the properties of the stellar populations of the galaxies in AC\\,118 by imposing that their evolution at $z \\sim 0$ mimics the overall distribution in the ($\\mathrm{K,V\\!-\\!K}$) plane for a local rich cluster.  With the aim of constraining the galaxy evolution scenarios, several studies have adopted a purely parametric approach, by comparing the observed properties in the CM diagram with those predicted by models that are based on different sets of parameters and that explore different assumptions on the probability distributions of such parameters. These studies also assume that the scatter in the CM diagram arises merely from age (but see Ferreras, Charlot, \\& Silk 1999).  Although the first epoch of star formation for the cluster early-type population seems to be constrained to high redshifts for almost all such models, further properties, as the last epoch and the spread of star formation activity, are more model dependent.  The procedure we adopted describes each galaxy of AC\\,118 by a stellar population model, which is constrained, when evolved to $z\\sim0$, to be bounded by the red sequence locus or by the region of blue galaxies of a rich nearby cluster (see Section 3.1). This is achieved by a suitable procedure which scales the magnitudes of the galaxy templates. We find that the best fitting models of AC\\,118 galaxies are able to match both the distributions in the ($\\mathrm{K,R\\!-\\!K}$) and ($\\mathrm{R\\!-\\!K}$,$\\mathrm{V\\!-\\!I}$) planes at $z\\sim0.3$, and the properties of the ($\\mathrm{K,V\\!-\\!K}$) color-magnitude distribution at $z \\sim 0$, i.e. slope and intrinsic scatter of the CM sequence, fraction of blue galaxies and luminosity function.  It is important to notice that such a result is not implicit in the method we used to scale the template magnitudes (see Section~\\ref{FITMOD}).  The constraint at $z \\sim 0$ largely reduces the region of input parameters available to the model. In particular, a sharp sequence arises in the metallicity-luminosity diagram (cfr. lower left panels of Figures~\\ref{CPLOT1} and ~\\ref{CPLOT2}), for which brighter galaxies are described by higher values of $Z$. The slope of the sequence is in full agreement with that derived by KA97 in the framework of the monolithic collapse/galactic wind model. It is interesting to notice that if we adopt the luminosity--weighted mean stellar metallicity of the KA97 models, the zero-points of the relations also coincide. The main difference between the results of KA97 and those of the present work is that we do not obtain the metallicity sequence on the basis of a particular galaxy evolution scenario.  The present data seem to exclude significant variations of star formation history along the CMR, and therefore confirm a pure metallicity interpretation, in which the bulk of the populations formed at high redshift ($z \\sim 5$). These results, however, do not describe the properties of all the stellar populations in AC\\,118: we find that about $20\\%$ of the points of our model do not follow any metallicity--luminosity relation, but are characterized by higher values of Z and more prolonged star formation activity ($\\tau>4~\\mathrm{Gyr}$). Since these objects do not show peculiar photometric properties in the colors-magnitude space, this result could be the consequence of a residual age--metallicity degeneracy. However, other possibilities can be explored. For instance, the scatter of the CM relation at a given luminosity could be partly due to the fact that more recently assembled galaxies have higher metallicity than older systems of similar luminosity (see Ferreras et al. 1999).  To study the dispersion in the CM diagram at $z\\sim0.3$, we computed the epoch $t_{90\\%}$ at which galaxies completed to form $90\\%$ of their stars. While for $\\mathrm{K<17}$ the corresponding redshift is greater than $z=1$, at faintest magnitudes ($\\mathrm{K}\\sim K^{\\star}+3$) we find that some galaxies ceased to form stars at epochs as low as $z\\sim0.5$. These results are in agreement with the general picture that fainter galaxies were more recently accreted from the field to the cluster environment and therefore ceased to form stars at later epochs (see KB01 for a wide discussion).  One half of the luminous mass present at $z\\sim0.3$ formed at $z>2$, and star formation continued at $z<1$ for $\\sim 20\\%$ of the stars. This result changes if we are neglecting the effect of the dust obscuration in a significant fraction of cluster galaxies. To investigate this subject, we adopted a simple model in which all the galaxies redder than the CM envelope at $z\\sim0.3$ are obscured by a uniform screen of dust.  The introduction of this model is also supported by the presence of few galaxies of AC\\,118 whose R-K color is too red with respect to the CMR. These objects may be accounted for as dusty galaxies with extensive on-going star formation activity (cfr. KB01). In the model with dust, the fraction of mass that forms at $z<1$ increases from $\\sim 20 \\%$ to $\\sim 30 \\%$."
    },
    "0206/astro-ph0206290_arXiv.txt": {
        "abstract": "The {\\it Extreme Ultraviolet Explorer\\/} ({\\it EUVE\\/}) satellite was employed for 5.46 days beginning on 1999 February 9.03 UT to acquire phase-resolved EUV photometric and spectroscopic observations of the AM~Her-type cataclysmic variable V834 Centauri. The resulting data are superior to those obtained by {\\it EUVE\\/} beginning on 1993 May 28.14 UT because the source was approximately three times brighter, the observation was four times longer and dithered, and {\\it ASCA\\/} observed the source simultaneously. Although we do not understand the EUV light curves in detail, they are explained qualitatively by a simple model of accretion from a ballistic stream along the field lines of a tilted ($[\\beta , \\psi ]\\approx [10^\\circ , 40^\\circ ]$) magnetic dipole centered on the white dwarf. In 1993 when the EUV flux was lower, accretion was primarily along the $\\varphi\\approx\\psi\\approx 40^\\circ $ field line, whereas in 1999 when the EUV flux was higher, accretion took place over a broad range of azimuths extending from $\\varphi\\approx \\psi\\approx 40^\\circ $ to $\\varphi\\approx 76^\\circ $. These changes in the accretion geometry could be caused by an increase in the mass-accretion rate and/or the clumpiness of the flow. The 75--140~\\AA \\ {\\it EUVE\\/} spectra are well described by either a blackbody or a pure-H stellar atmosphere absorbed by a neutral hydrogen column density, but constraints on the size of the EUV emission region and its UV brightness favor the blackbody interpretation. The mean 1999 EUV spectrum is best fit by an absorbed blackbody with temperature $kT\\approx 17.6$ eV, hydrogen column density $N_{\\rm H}\\approx 7.4 \\times 10^{19}~\\rm cm^{-2}$, fractional emitting area $f\\approx 10^{-3}$, 70--140~\\AA \\ $\\rm flux\\approx 3.0\\times 10^{-11}~\\rm erg~cm^{-2}~s^{-1}$, and luminosity $L_{\\rm soft}\\approx 7.2\\times 10^{32}\\, (d/100~{\\rm pc})^2~\\rm erg~s^{-1}$. The ratio of the EUV to X-ray luminosities is $L_{\\rm soft} /L_{\\rm hard}\\approx 40$, signaling that some mechanism other than irradiation (e.g., blob heating)  dominates energy input into the accretion spot. The 1999 SW hardness  ratio variation can be explained by minor variations in $kT$ and/or $N_{\\rm H}$, but instead of tracking the SW count rate variation, the hardness ratio variation was sinusoidal, with a minimum (maximum) when the accretion spot was on the near (far) side of the white dwarf, consistent with the trend expected for an atmosphere with an inverted temperature distribution. ",
        "introduction": "Polars or AM~Her stars are a class of semidetached binaries composed of a low-mass main-sequence secondary and a strongly magnetic ($B\\approx 10$--100 MG) white dwarf primary. The strong field locks the white dwarf into corotation and the accreting matter is channeled along the field lines for much of its trajectory from the secondary's inner Lagrange point to a spot near the white dwarf magnetic pole. To match boundary conditions, the flow passes through a strong shock far enough above the star for the hot ($kT\\le kT_{\\rm shock}= 3G\\Mwd \\mu m_{\\rm H}/8\\Rwd \\sim 20$ keV), post-shock matter to cool and come to rest at the stellar surface. This plasma cools via bremsstrahlung and line emission in the X-ray bandpass and cyclotron emission in the optical and near-IR. Roughly half of this radiation is emitted outward and is observed directly; the other half is emitted inward, where it is either reflected or absorbed by the white dwarf surface. In addition to  radiative heating, the white dwarf surface can be heated by blobs of material which penetrate to large optical depths before thermalizing  their kinetic energy. For a mass-accretion rate $\\Mdot =10^{16}~\\rm g~s^{-1}$ and a relative spot size $f=10^{-3}$, the white dwarf surface is heated to a temperature $kT=k(G\\Mwd\\Mdot /4\\pi\\sigma f\\Rwd ^3)^{1/4}\\sim 20$~eV, hence produces a radiation spectrum which peaks in the EUV.  Recent progress in our understanding of the accretion spots of polars has come from photometric and spectroscopic observations with {\\it EUVE\\/}. The state of our understanding of the {\\it EUVE\\/} spectra of polars as of mid-1998 is described by \\citet{mau99}, who fit blackbody, pure-H stellar atmosphere, and solar abundance stellar atmosphere models to the phase-average {\\it EUVE\\/} spectra of nine polars with useful data in the archive. Among the models tested, the blackbody parameterization gave the best fits to the data, with blackbody temperatures $kT_{\\rm bb}\\approx 15$--25~eV and hydrogen column densities $\\log N_{\\rm H}({\\rm cm^{-2}})\\approx 19$--20. Wishing to increase the number of polars with good signal-to-noise ratio EUV spectra, in 1998 we obtained approval for a 130 ks {\\it EUVE\\/} observation of V834~Cen, whose 40 ks archival spectrum had a {\\it peak\\/} signal-to-noise ratio of only eight in 0.54~\\AA \\ bins. Coincidentally, in late 1998 M.~Ishida obtained approval for an additional {\\it ASCA\\/} X-ray observation of V834~Cen, so it was arranged that these observations should be obtained simultaneously. A description of the X-ray light curves and spectra from the {\\it ASCA\\/} observation is provided by \\citet{ish99} and \\citet{ter01}. Below we present the analysis of the EUV light curves and spectra from both the original 1993 and the new 1999 {\\it EUVE\\/} observations of V834 Cen.  V834 Cen (nee E1405$-$451) is a well-studied polar with an orbital period $P_{\\rm orb}=101.5$ minutes \\citep{mas83} and the high ($V\\approx14$) and low luminosity states ($V\\approx 17$) typical of this class of binaries. The magnetic nature of the white dwarf is demonstrated directly by the linear and circular polarization present in its high states \\citep{cro86, cro89} and the Zeeman absorption features and cyclotron emission features present in its low states \\citep{sch90, fer92}, from which a magnetic field strength of 23~MG is inferred. The distance $d$, white dwarf mass $\\Mwd $, binary inclination $i$, and the colatitude $\\beta $ and azimuth $\\psi $ of the accretion spot are all uncertain at some level, but \\citet{puc90} find $d>77$ pc based on a detection of the secondary in a near-IR spectrum; \\citet{cro99} find $\\Mwd=0.54$--$0.64~\\Msun $ based on fits of a {\\it Ginga\\/} X-ray spectrum, \\citet{ram00} finds $\\Mwd = 0.64$--$0.68~\\Msun $ based on fits of an {\\it RXTE\\/} X-ray spectrum, and \\citet{sch93} use phase-resolved optical spectra to determine $\\Mwd =0.66^{+0.19}_{-0.16}~\\Msun $ for $i={50^\\circ } \\, ^{+10^\\circ }_{-5^\\circ }$; \\citet{cro88} advocates $i=45^\\circ \\pm 9^\\circ $, $\\beta =25^\\circ\\pm 5^\\circ $, and $\\psi = 40^\\circ \\pm 5^\\circ $ based on values compiled from the literature.  The plan of this paper is as follows. In \\S 2 we describe the {\\it EUVE\\/} observations, in \\S 3 we present the {\\it EUVE\\/} deep survey (DS) photometer and {\\it ASCA\\/} count rate light curves, in \\S 4 we present the {\\it EUVE\\/} short wavelength (SW) spectrometer count rate and hardness ratio light curves, in \\S 5 we present the mean and phase-resolved SW spectra, in \\S 6 we provide a discussion and interpretation of these data, and in \\S 7 we close with a summary of our results. ",
        "conclusions": "In the previous sections we have described the phenomenology of the EUV light curves and spectra of V834 Cen measured by {\\it EUVE\\/} in 1993 May and 1999 February. To help understand these data, we note the following, based on our general understanding of polars and the results from optical studies of V834 Cen by \\citet{ros87}, \\citet{cro88, cro89}, and \\citet{sch93}. (1) The binary inclination $i\\approx 45^\\circ\\pm 9^\\circ $ and the accretion spot colatitude and azimuth are respectively $\\beta\\approx 25^\\circ\\pm 5^\\circ $ and $\\psi\\approx 40^\\circ\\pm 5^\\circ $. (2) In V834 Cen only one accretion spot is visible for all orbital phases because the sum of the spot colatitude and binary inclination is less than $90^\\circ $. (3) Because the binary inclination is greater than the spot colatitude, the accretion stream passes through the line of sight to the accretion region once per binary revolution. (4) The narrow and broad components of the optical emission lines of polars are understood to be due to, respectively, the heated face of the secondary and the base of the accretion stream. (5) Blue-to-red zero crossing of the narrow components of optical emission lines occurs at binary phase $\\phi\\approx 0$, placing the secondary (white dwarf) on the near (far) side of the binary at that phase. (6) Maximum blueshift of the broad components of optical emission lines occurs at $\\phi\\approx 0.42$, so at that phase the accretion stream points most directly toward us. (7) The linear polarization spike also occurs at $\\phi\\approx 0.42$, so at that phase the accretion column is on the plane of the sky. (8) We expect that the accretion column will point most directly toward us approximately $180^\\circ $ later, at $\\phi\\sim 0.92$. The circular polarization, EUV, and X-ray light curves are all eclipsed near this phase because of the passage through the line of sight of the (pre-shock) accretion stream and (post-shock) accretion column.  To illustrate these results, we constructed the graphic shown in the left panel of Figure~8, which includes the nominal position on the white dwarf of the accretion spot ({\\it filled trapezoid\\/}); the 1993 and 1999 DS light curves ({\\it grey and black polar histograms, respectively\\/}); and the phases of the maximum blueshift of the broad component of optical emission lines, the spike in linear polarization light curves, and the dip in circular polarization light curves. It is seen that the primary eclipse of the EUV light curves occurs when the accretion spot points toward the observer, while the dip in circular polarization light curves appears to occur {\\it between\\/} the primary and secondary eclipses of the EUV light curves.  To help envision the accretion geometry of V834 Cen, we constructed a simple model, shown in the right panels of Figure~8, of the path of material from the secondary's inner Lagrange point to the white dwarf surface in the vicinity of the upper magnetic pole. The model consists of the ballistic stream for a nonmagnetic semidetached binary and the field lines of a tilted magnetic dipole centered on the white dwarf for an orbital period $P_{\\rm orb}= 101.5$ minutes, mass ratio $q=5.08$ (determining the trajectory of the ballistic stream), white dwarf mass $\\Mwd = 0.66~\\Msun $ (determining the white dwarf radius $\\Rwd =8.1 \\times 10^8$ cm), magnetic colatitude $\\beta = 10^\\circ $, and magnetic azimuth $\\psi = 40^\\circ $. Field lines are drawn for azimuthal angles $\\varphi = 0^\\circ , 10^\\circ , 20^\\circ , \\ldots , \\psi +90^\\circ $. If material in the ballistic stream makes it beyond $\\varphi\\approx\\psi +90^\\circ\\approx 130^\\circ $, it will accrete preferentially onto the lower magnetic pole, hence disappears from consideration. In this simple model, the material lost by the secondary travels along the ballistic stream until it is threaded by the white dwarf magnetic field; it then leaves the orbital plane and follows the magnetic field lines down to the white dwarf surface in the vicinity of the magnetic poles. In reality, the trajectory of the accreting material will be affected by the magnetic field and the magnetic field will be distorted by the accreting material; our simple model is a first-order approximation of the path of the accreting material which ignores these complications.  In the right panels of Figure 8 we see that if material accretes onto the white dwarf from the full range of possible azimuthal angles, it will produce a long, thin, arc-shaped accretion region on the white dwarf surface at the footpoints of the magnetic field lines. Details of the accretion region are more apparent in Figure~9, which shows the white dwarf and magnetic field lines for a binary inclination $i=50^\\circ $ and binary phases $\\phi = 220^\\circ , 230^\\circ , 240^\\circ , \\ldots , 360^\\circ $ ($\\phi = 0.61$--1.0). For the assumed parameters, the accretion region extends about $90^\\circ $ in azimuth and has a variable offset of about $20^\\circ $ from the magnetic pole, determined by the varying distance (a maximum of $37.5~\\Rwd $ for $\\varphi =0^\\circ $ and a minimum of $5.9~\\Rwd $ for $\\varphi=130^\\circ $) to the ballistic stream from the white dwarf. Accretion along the $\\varphi =\\psi =40^\\circ $ field line (the thick  black curves in Figs.~8 and 9) will produce an accretion spot at $[\\beta , \\psi ] =[26^\\circ , 40^\\circ ]$, consistent with the values advocated by \\citet{cro88}; the other model parameters are consistent with the radial velocity solution of \\citet{sch93}.  \\subsection{EUV Light Curves}  Whether a spot or an arc, the flux at Earth from the accretion region will vary with binary phase because of the varying projected surface area of the accretion region and the phase-dependent obscuration by infalling material. Ignore for the moment obscuration and consider the simplest possible model of the EUV light curves of V834 Cen, that of a small spot ``painted'' on the white dwarf surface: its projected surface area hence light curve varies as $\\sin i\\,\\sin\\beta\\, \\cos\\, (\\phi +\\psi) + \\cos i\\,\\cos \\beta $. Such a simple model fails to reproduce even the upper envelope of the EUV light curves of V834 Cen, and although linear a superposition of such functions (i.e., multiple spots) can reproduce some aspects of the data, the fits are not unique. Consider next an empirical model, using the EUV light curve of AM~Her \\citep{pae96} as a template to explain the 1999 EUV light curve of V834 Cen. Aligning the two light curves on the respective linear polarization ephemerides, the AM~Her light curve reproduces well the V834 Cen light curve between binary phases $\\phi\\approx 0.0$--0.17. By shifting the AM~Her light curve by $\\Delta\\phi = 0.5$, it reproduces well the V834 Cen light curve between binary phases $\\phi\\approx 0.17$--0.6. From this exercise, one is led to imagine that the EUV emission region of V834 Cen is composed of two spots separated in azimuth by about $180^\\circ $: the first spot produces the falling portion of the V834 Cen light curve between $\\phi\\approx 0.0$--0.17, while the second spot produces the rising portion of the light curve between $\\phi\\approx 0.17$--0.6; the remaining portion of the light curve is affected by the primary and secondary eclipses. Independent support for two accretion spots is provided by \\citet{cro89}, who advocates two magnetic poles separated in azimuth by about $230^\\circ $ to explain the occasional presence of two waves in the positional angle of the linearly polarized optical flux of V834 Cen. Such a two-spot model allows us to understand two features of the EUV light curves of V834 Cen. First, the secondary minimum at $\\phi\\approx 0.17$ discussed in \\S 3 can be understood as being the interface between the falling portion of the light curve of the leading spot and the rising portion of the trailing spot. Second, the change between 1993 and 1999 in the relative strengths of the EUV light curves between phases $\\phi \\approx 0.0$--0.17 and $\\phi\\approx 0.17$--0.6 (cf.\\ the top panel of Fig.~1) can be understood if the leading (trailing) accretion spot was brighter in 1993 (1999). Unfortunately, this model predicts that the accretion spots lie at $\\phi\\approx 0.05$ and $\\phi\\approx 0.55$, which appears not to be the case. These considerations teach us the folly of attempting to understand in detail the EUV light curves of V834 Cen (or polars in general, see \\citealt{sir98}).  Returning to the discussion of the previous section, we expect that if the ballistic stream is threaded by the white dwarf magnetic field over a {\\it narrow\\/} range of azimuths, the infalling material will form a narrow stream leading down to a small spot on the white dwarf surface, whereas if the ballistic stream is threaded by the magnetic field over a {\\it broad\\/} range of azimuths, the infalling material will form an ``accretion curtain'' (as it called in intermediate polars) leading down to an extended arc on the white dwarf surface. Narrow streams are implied by the optical eclipse studies of HU~Aqr and UZ~For by \\citet{har99} and \\citet{kub00}, respectively, but an extended accretion region is invoked by \\citet{cro89} in V834 Cen to explain changes in the shape of the optical light curve and the changing morphology of the linear polarization and position angle curves. We propose that the EUV light curves of V834 Cen are explained by accretion over a narrow range of azimuths in 1993 and a broad range of azimuths in 1999. Specifically, we propose that in 1993 accretion was primarily along the $\\varphi\\approx\\psi\\approx 40^\\circ $ field line, resulting in a narrow eclipse of a small accretion spot, while in 1999 accretion was more spread out in azimuth: the primary accretion channel remained the $\\varphi\\approx\\psi\\approx40^\\circ $ field line, but field lines $\\approx 36^\\circ $ further along the ballistic stream became active, resulting in the secondary eclipse observed at $\\phi\\approx 0.79$. Happily, the proposed shift of the accretion rate further along the ballistic stream qualitatively explains the reduction in the relative strength of the EUV light curve between $\\phi\\approx 0.0$--0.17. What caused the changes in the accretion geometry between 1993 and 1999? One option is a higher accretion rate, since in 1999 the EUV flux was approximately three times higher than in 1993. Another option is enhanced clumping of the flow, since greater clumping will result in magnetic threading over a broader range of azimuths {\\it and\\/} more efficient heating of the accretion spot as blobs of material crash into the stellar surface.  \\subsection{EUV Spectra}  Although we have appealed to obscuration by the pre- and post-shock accretion flow to explain the primary and secondary eclipses in the EUV light curves of V834 Cen, Figure~3 makes clear that this obscuration cannot be due to photoelectric absorption. First, the amplitude of the SW hardness ratio variation is too small. Assuming that the mean EUV spectrum of V834 Cen is an absorbed blackbody with $kT=17.6$ eV and $N_{\\rm H}=7.4\\times 10^{19}~\\rm cm^{-2}$ (cf.\\ Table~1 and Fig.~7), the SW hardness ratio equals 1.1, as observed. Doubling $N_{\\rm H}$ decreases the SW count rate by a factor of $\\approx 10$, but the SW hardness ratio increases by a factor of $\\approx 4$. In contrast, during the 1999 observation, the SW count rate varied by a factor of $\\gax 10$ while the hardness ratio varied by only $\\approx 20$\\%. Similar results apply even if H and He are singly ionized, and only the \\ion{He}{2} bound-free opacity is left to absorb the EUV flux from the accretion spot. It is actually quite easy to remove the photoelectric opacity of all these ions: in a collisionally ionized plasma, H is fully ionized for $T\\gax 25$ kK and He is fully ionized for $T\\gax 100$ kK. Above that temperature, only Thomson opacity is available to reduce the EUV flux, and an electron column density $N_{\\rm e}= -\\ln(0.1)/\\sigma _{\\rm T} = 3.5\\times 10^{24}~\\rm cm^{-2}$ is required to reduce the EUV flux by a factor of ten. Second, the SW hardness ratio variation does not track the SW count rate variation. Instead, at least during the 1999 observation (see Fig.~3), the hardness ratio variation was sinusoidal, with a minimum (maximum) when the accretion spot was on the near (far) side of the white dwarf. Evidently, as the accretion spot rotates to the far side of the white dwarf, its spectrum systematically hardens. Figure~7 shows that the detailed spectral fits indicate that this change is best explained by decreasing $kT$ {\\it and\\/} increasing $N_{\\rm H}$, but within the 68\\% contours this can be accomplished simply by increasing $N_{\\rm H}$, and within the 90\\% contours this can be accomplished simply by increasing $kT$. The last option implies that the accretion region is harder when observed at shallow viewing angles. This is the trend expected for an accretion spot subjected to irradiation by the million-degree plasma in the accretion column \\citep{tee94}, although the affect is predicted to be small in the 75--140~\\AA \\ bandpass and the large ratio of hard to soft X-ray luminosities (see below) argues that irradiation does not dominate the energy deposition into the accretion spot.  The mean 1993 and 1999 and the phase-resolved 1999 {\\it EUVE\\/} spectra of V834~Cen can be fit by a blackbody or a pure-H stellar atmosphere model absorbed by a neutral hydrogen column density with the parameters listed in Table~1. Compared to the blackbody models, the stellar atmosphere models are cooler, larger, and more luminous. In fact, as shown by Figures~5 and 7, the stellar atmosphere models are {\\it too\\/} big and {\\it too\\/} luminous (too bright in the UV) to explain the observations of V834 Cen. First, the stellar atmosphere model fails to account for the EUV spectrum acquired during the hard phase of the 1999 observation because the required angular size is larger than that of the white dwarf. Second, the stellar atmosphere model fails to account for the mean 1999 EUV spectrum because it produces too much flux in the UV unless the reddening $A_V\\ge 1$. Similar problems are met applying this model to other polars \\citep{mau99}, so we conclude that the pure-H stellar atmosphere model cannot in general explain the EUV of these magnetic CVs.  Finally, having {\\it simultaneously\\/} measured the EUV and X-ray spectra of V834 Cen in 1999, we are in a position to determine the ratio of the luminosities of the accretion column and accretion spot. Table~1 shows that the absorbed blackbody fit to the mean 1999 EUV spectrum yields best-fit parameters $kT\\approx 17.6$ eV, $N_{\\rm H} \\approx 7.4\\times 10^{19}~\\rm cm^{-2}$, fractional emitting area $f\\approx 10^{-3}$, 70--140~\\AA \\ $\\rm flux\\approx 3.0\\times 10^{-11}~\\rm erg~cm^{-2}~s^{-1}$, and luminosity $L_{\\rm soft}\\approx 7.2\\times 10^{32}\\, (d/100~{\\rm pc})^2~\\rm erg~s^{-1}$. The averaged 2--10 keV flux measured by {\\it ASCA\\/} was $1.5\\times 10^{-11}~\\rm erg~cm^{-2}~s^{-1}$ \\citep{ish99}, implying $L_{\\rm hard}\\approx 1.8\\times 10^{31}\\, (d/100~{\\rm pc})^2~\\rm erg~s^{-1}$, hence $L_{\\rm soft}/L_{\\rm hard}\\approx 40$ (to within a factor of $\\approx 2$). This imbalance between the soft accretion spot and hard accretion column luminosities is the famous ``soft X-ray problem,'' and signals that some mechanism other than irradiation (e.g., blob heating) dominates energy deposition into the accretion spot."
    },
    "0206/nucl-th0206009.txt": {
        "abstract": "We study the hadron-quark phase transition in the interior of neutron stars (NS's). We calculate the equation of state (EOS) of hadronic matter using the Brueckner-Bethe-Goldstone formalism with realistic two-body and three-body forces, as well as a relativistic mean field model.  For  quark matter  we  employ  the  MIT  bag model constraining the bag constant by using the indications coming from  the recent experimental  results obtained  at  the CERN  SPS on  the formation of a quark-gluon plasma. We find necessary to introduce a density dependent bag parameter, and the corresponding consistent thermodynamical formalism. We calculate the structure of NS  interiors  with  the EOS comprising both phases, and we find that the NS maximum masses fall in a relatively  narrow  interval, $1.4\\,M_\\odot  \\leq M_{\\rm  max}  \\leq  1.7\\,M_\\odot$. The precise value of the maximum mass turns out to be only weakly correlated with the value of the energy density at the assumed transition point in nearly symmetric nuclear matter. ",
        "introduction": "The  properties of nuclear matter  at high density play a crucial role for building models of neutron stars \\cite{shapiro}. The  observed NS's masses are in the range of $\\approx (1-2)  M_\\odot$  (where  $M_\\odot$  is  the  mass of the sun, $M_\\odot = 1.99\\times 10^{33}$g), and the radii are of the order of 10 km. The characteristics of the core of the NS's influences most strongly the value of the  maximum mass. The matter inside this core possesses densities ranging from a few times $\\rho_0$ $(\\approx  0.17\\;{\\rm  fm}^{-3}$,  the normal nuclear matter density) to one order of magnitude higher. Moreover, the equation of state at such high densities is the main ingredient to determine the structure parameters of NS's, such as mass and radius. Therefore, a detailed  knowledge  of  the EOS  is required for densities $\\rho \\gg \\rho_0$, where a description of matter only in terms of nucleons and leptons may be inadequate. In fact, at densities $\\rho \\gg \\rho_0$ several species of other particles,  such as  hyperons  and  $\\Delta$ isobars, may appear, and meson condensations may take place; also, ultimately, at very high densities, nuclear matter is expected to undergo a phase transition to a quark-gluon plasma \\cite{quark}. The specific goal of the theory is to study the nature of this plasma and understand the phase transitions between different states. However, the exact value of the transition density to  quark  matter  is unknown and still a matter of recent debate not only in astrophysics, but also within the theory of high energy heavy ion collisions.  In this paper, we propose a method to determine a range of values of the maximum mass of NS's taking  into  account the phase transition from hadronic matter to quark matter inside the neutron star. The transition point is constrained from recent heavy-ion collision data. Therefore to perform such calculations, we describe the hadron phase of matter by using two different equations of state, {\\rm i.e.}, a microscopic non-relativistic EOS obtained in the Brueckner-Bethe-Goldstone (BBG) theory \\cite{book}, and  a more phenomenological relativistic mean field model \\cite{serot}. The deconfined quark matter phase is treated by adopting the popular MIT bag model \\cite{chodos}. In a previous paper \\cite{bbsss} the bag ``constant'', $B$, which  is  a  parameter  of  the  bag  model, was constrained  to  be  compatible  with  the  recent  experimental results obtained at CERN on the formation of a quark-gluon plasma \\cite{heinz}, recently confirmed by RHIC preliminary results \\cite{blaizot}.  However, it is not obvious if the information on the nuclear EOS from high energy heavy ion collisions can be related to the physics of neutron star interiors. The possible quark-gluon plasma produced in heavy ion collisions is expected to be characterized by small baryon density and high temperature, while the possible quark phase in neutron stars appears at high baryon density and low temperature. If one adopts for the hadronic phase a non-interacting gas model of nucleons, antinucleons, and pions, the original MIT bag model  predicts that the deconfined phase occurs at an almost constant value of the quark-gluon energy density, {irrespective} of the thermodynamical conditions of the system \\cite{gavai}. For this reason, it is popular to draw the transition line between the hadronic and the quark phase at a constant value of the energy density, which was estimated to fall in the interval between 0.5 and 2 GeV fm$^{-3}$ \\cite{mul}. This is consistent with the value of about 1 GeV fm$^{-3}$ reported by CERN experiments. The close relation between the physics of neutron stars and of heavy ion collisions is also emphasized by a recent conjecture that there could be three phases in heavy-ion collisions at SPS and RHIC energies, equivalent to the pure quark phase, mixed phase, and pure hadron phase appearing in neutron stars. These three phases correspond to (a) an explosive hard quark-gluon phase, (b) a mixed soft phase (a sort of plateau), and (c) a hadronic phase. Considering these three phases in a heavy-ion collision model, the first available RHIC data could be well described \\cite{shuryak}.  The value of 1 GeV fm$^{-3}$ must be considered only an indicative estimate of the transition energy density at zero or nearly zero temperature, as needed in neutron star studies, and it appears mandatory to explore the sensitivity of the results on the precise value of the assumed transition energy density. In this work we present systematic calculations of neutron star structure, where the hadronic EOS, which can be considered well established, is implemented with the possible transition to the deconfined phase described by different parametrizations of the MIT bag model. The transition energy density in nearly symmetric nuclear matter at zero temperature is allowed to vary within a range of values which can be considered still compatible with the CERN and RHIC data. The calculations will indicate the sensitivity of the results on the assumed transition point and the possible correlation between neutron star properties and the transition energy density value. In particular, we will see that the maximum neutron star mass is only weakly correlated with the transition energy density value.  This paper is organized as follows. In Sec.~II we discuss the EOS for the hadronic phase of a neutron star, {\\rm i.e.}, the BBG and the relativistic mean field models. In Sec.~III we apply the MIT bag model to the description of the quark phase of the neutron star. In Sec.~IV we present our results and finally in Sec.~V we draw some conclusions. ",
        "conclusions": "We studied neutron star properties, in particular NS's maximum masses, using  an  EOS  which combines reliable EOS's for hadronic matter and a bag model EOS for quark matter.  We found that a density dependent $B$ is necessary to get the transition to the quark-gluon plasma in nearly symmetric nuclear matter at an energy density which is well above saturation density and in a range of values which can be considered compatible with the CERN-SPS and RHIC findings on  the  phase  transition from  hadronic  matter  to  quark matter. We considered a wide range of values, from 0.8 $\\rm GeV~fm^{-3}$ to 1.5 $\\rm GeV~fm^{-3}$, in order to establish the sensitivity of the results on the assumed value of the transition energy density. \\par For a given value of the transition density for symmetric nuclear matter, the corresponding transition in neutron star matter, i.e., beta stable matter, occurs in general at substantially lower energy density. It is essential, in this respect, that in the calculations strange matter is included and allowed to develop inside neutron star matter, since the appearence of strange matter tends in general to soften the EOS. The results show that the NS maximum mass is clearly correlated with the assumed value of the transition energy density. For a given transition density, the maximum mass falls in a narrow range, nearly independent of the details of the parametrization of the bag model. As the transition density is made to vary, the value of the maximum mass is shifted. In general it decreases at increasing value of the transition energy density if the hadron EOS is computed within the microscopic BHF scheme. The trend is reversed with the hadron EOS computed within the relativistic mean field method. However, this correlation appears to be rather weak, and the full range of possible values of the maximum mass turns out to be  between 1.4 and 1.7 solar masses.  The value of the maximum mass is mainly determined by the quark component of the neutron star and by the corresponding EOS. In this sense, one can say that the value of the neutron star maximum mass can be a good testing ground for the quark EOS, rather than the hadron EOS. Indeed, the value of the maximum mass of neutron stars obtained according to our analysis appears robust with respect to the uncertainties of the nuclear EOS, and the obtained range of values is mainly due to the uncertainties of the quark EOS.\\par Other recent calculations of neutron star properties employing various RMF nuclear EOS's together with either effective mass bag model  \\cite{bag}  or Nambu-Jona-Lasinio  model  \\cite{njl}  EOS's  for quark matter, also give maximum masses of only about $1.7\\,M_\\odot$, even though not constrained by hints coming from the CERN-SPS and RHIC data. Therefore, according to our results, the experimental observation of a heavy ($M > 1.8 M_\\odot$) neutron star, as claimed recently by some groups \\cite{kaaret}($M \\approx 2.2 M_\\odot$), if confirmed, would suggest that serious drawbacks are present for the possible description of the high-density phase of quark matter within the bag model.  \\newpage"
    },
    "0206/hep-ph0206210_arXiv.txt": {
        "abstract": " ",
        "introduction": "Our observable universe is four-dimensional. However fundamental theories, like for example string theories, require for consistency more than four space-time dimensions.  In the standard Kaluza-Klein approach \\cite{kk} the extra dimensions are compact and curled up to very small size, roughly of the order of the Planck scale, ${\\ell }_P\\sim 10^{-33} {\\rm cm}$, implying that space-time is effectively four-dimensional. However recently, with the discovery of brane solutions to string/M-theory, the idea of large extra dimensions has attracted much attention \\cite{hw}.  In this last approach, it is proposed that ordinary matter is confined to a three-dimensional sub-manifold, a three-brane, which is embedded in a higher dimensional space-time; the graviton on the other hand is allowed to propagate freely through the whole space-time. Neglecting the brane tension, i.e. the energy density per unit three-volume of the brane, and considering compact dimensions, one re-introduces the Kaluza-Klein picture.  However, in this case the extra dimensions do not have to be small.  Newtonian gravity has been tested and found to hold down to scales of order of 1 mm \\cite{newgra}. Below this scale, gravity could in principle be higher than four-dimensional. In higher dimensional gravity, the four-dimensional Planck scale is no longer a fundamental scale; the higher dimensional Planck scale, $M_{\\rm D}$, becomes instead the fundamental scale. This allows to explain the huge difference between the electroweak and the four-dimensional Planck scale, known as the hierarchy problem. Assuming for simplicity that the $n$ extra dimensions form an $n$-torus which has the same radius $R$ in each direction, and using Gauss' law, the $D$-dimensional and the four-dimensional Planck masses, $M_4$ and $M_{\\rm D}$ respectively, are related by \\cite{ADD} \\cite{ruba} \\begin{eqnarray} M_4^2=R^nM_{\\rm D}^{n+2}~, \\end{eqnarray} where $R^n$ is the volume of the $n$ extra dimensions. Thus, taking $M_{\\rm D}$ to be of the order of the electroweak scale, $M_{\\rm D}\\sim 1 {\\rm TeV}$, the huge difference between $M_4$ and the electroweak scale can be explained as due to the large size of the extra dimensions.  Here we are going to study the effects of time-varying extra dimensions on the graviton production.  We assume that the background space-time can be written as a direct product of an internal space, i.e. the extra dimensions, and a four-dimensional external space-time. Furthermore, during an initial phase, the internal space is contracting and the four-dimensional external space-time is expanding. At some time this is matched to a radiation dominated flat Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) universe with the internal dimensions frozen at constant size.  Due to the changing background metric, there will be particle production. Here we will be in particular concerned with the production of gravitons. Using observational bounds on graviton spectra it is possible to derive a relation between the $D$-dimensional Planck mass and the temperature at the time of transition.  Usually it is assumed that there are no massive modes excited, i.e. momenta lying in the internal space are not taken into account.  Here however, we are going to assume that the internal momenta are excited and we are interested in their effect on the spectral energy density in four dimensions.  In order to determine this effect, the internal momenta are integrated out.  This will lead to a final expression for the spectral energy density, which will be just depending on the four-dimensional momenta.  The formalism to deal with metric perturbations has already been developed in Ref.~\\cite{abe}. Particle production in higher dimensional space-times has been also the subject of Refs. \\cite{GV} and \\cite{giov}. ",
        "conclusions": "We have discussed graviton production in the context of a higher dimensional model, which from a multi-dimensional phase where the extra dimensions are contracting and the external dimensions are expanding, enters into an effectively four-dimensional universe with static extra dimensions.  The momenta in the extra space were taken into account.  To find an estimate of their contribution to the energy density in the four-dimensional space-time, the momenta in the internal space were integrated out.  The contributions of the internal momenta tend to dominate the gravitational wave spectrum, up to the point of forcing the universe to recollapse at an early stage.  However, one can assume that there is an upper cut-off for the internal momenta, given by the maximal frequency of the four-dimensional (external) space-time.  The gravitational wave spectrum is constrained by observations.  With the assumption that the standard relation between the Hubble parameter and the temperature holds up to the transition time, we derived an upper bound on the temperature at the beginning of the radiation dominated era.  Using that Newtonian gravity has been tested sucessfully down to scales of the order of 1 mm, one obtains a lower bound on the temperature at the transition scale, which depends on the number of extra dimensions.  This lower bound on the reheating temperature is, in general, lower than the reheating temperature of standard inflationary models. However, it can be raised if one assumes smaller extra dimensions, which lead to a higher $D$-dimensional Planck mass. Our final conclusion is that the size $R$ of the internal space must be much larger than Planck length $\\ell_{\\rm Pl}$, namely $R\\gg \\ell_{\\rm Pl}$."
    },
    "0206/hep-ph0206026_arXiv.txt": {
        "abstract": "We study the cosmic density perturbations induced from fluctuation of the amplitude of late-decaying scalar condensations (called $\\phi$) in the scenario where the scalar field $\\phi$ once dominates the universe.  In such a scenario, the cosmic microwave background (CMB) radiation originates to decay products of the scalar condensation and hence its anisotropy is affected by the fluctuation of $\\phi$.  It is shown that the present cosmic density perturbations can be dominantly induced from the primordial fluctuation of $\\phi$, not from the fluctuation of the inflaton field.  This scenario may change constraints on the source of the density perturbations, like inflation.  In addition, a correlated mixture of adiabatic and isocurvature perturbations may arise in such a scenario; possible signals in the CMB power spectrum are discussed.  We also show that the simplest scenario of generating the cosmic density perturbations only from the primordial fluctuation of $\\phi$ (i.e., so-called ``curvaton'' scenario) is severely constrained by the current measurements of the CMB angular power spectrum if correlated mixture of the adiabatic and isocurvature perturbations are generated. ",
        "introduction": "\\label{sec:introduction} \\setcounter{equation}{0}  In the recent years, observation of the cosmic microwave background (CMB) anisotropy has been greatly improved.  After the discovery of the CMB anisotropy by COBE \\cite{APJ464L1} at the angular scale $\\theta\\gtrsim 7^\\circ$, there have been many efforts to improve the measurements.  In particular, recent balloon and ground-based experiments observed the CMB anisotropy at smaller angular scale of $\\theta\\sim O(0.1^\\circ)$ \\cite{aph0104460,aph0104459,aph0104489}. Then, the observation of the CMB anisotropy is expected to be greatly improved by the on-going and future satellite experiments, MAP \\cite{MAP} and PLANCK \\cite{PLANCK}; after these experiments, the CMB angular power spectrum $C_l$ will be determined at $O(1\\ \\%)$ level up to the multipole $l\\lesssim 1000-2000$.  With these measurements, our understanding of the evolution of the universe is also being improved.  Importantly, since the CMB power spectrum is sensitive to origins and evolutions of the cosmic density perturbations (as well as to the cosmological parameters), we are now able to constrain scenarios of generating the cosmic density perturbations.  Among various scenarios, inflation \\cite{PRD23-347} is probably the most popular and well-motivated one to provide the source of the cosmic density perturbations.  In the simplest case, all the components in the universe (other than the cosmological constant), like the photon, baryon, cold dark matter (CDM), and so on, originate to decay products of the inflaton field $\\chi$ which is the scalar field responsible for the inflation.  The inflaton field fluctuates during the inflation and it becomes the source of the cosmic density perturbations.  One of the most important consequences of such a scenario is that, since all the components in the universe are produced from the inflaton field, there is no entropy perturbation between any of two components and hence the density fluctuations become adiabatic.  Importantly, assuming the standard evolution of the universe after the inflation and using a reasonable set of the cosmological parameters, it is now widely believed that the observed CMB power spectrum is consistent with the one predicted from the scale-invariant spectrum which can arise from some class of inflation models \\cite{hep-ph_0201264}.  From the particle-physics point of view, however, the simplest scenario may not be the case and there exist possible sources of the cosmic density perturbations other than the inflaton field.  In particular, in various scenarios, light scalar fields are introduced which dominate the universe at early epochs.  (We denote such a scalar field as $\\phi$.) For example, in the scenario of Affleck-Dine baryogenesis \\cite{NPB249-361}, condensation of the squark and slepton fields is converted to the baryon-number asymmetry of the universe. If this is the case, there may exist an epoch when the universe is dominated by the Affleck-Dine field.  Then, at the time of the decay of the Affleck-Dine field, the late-time entropy production occurs. Other candidate is flat directions in the theory space.  In particular, in the superstring theory \\cite{Polchinski}, there are various flat directions parameterized by scalar fields called moduli fields.  Such moduli fields are expected to acquire masses from the effect of the supersymmetry breaking and hence their masses can be much lighter than the Planck scale.  If the initial amplitudes of the moduli fields are large, the universe is once dominated by the moduli field and is reheated at the time of the decay of the moduli fields. (See Refs.\\ \\cite{PRL131-59,heavy-moduli} for the scenario with cosmological moduli fields.)  In addition, in Refs.\\ \\cite{snu-leptogen} it is pointed out that the baryon-number asymmetry of the universe can be explained if the universe was once dominated by the condensation of the right-handed scalar neutrino.  In this case, the universe also experiences the late-time entropy production.  In addition, axion-like scalar field is proposed as a seed of the cosmic density perturbations \\cite{axion=seeds} in the pre-big-bang \\cite{PreBigBang} and the ekpyrotic \\cite{ekpyrotic} scenarios.  Such scalar fields may acquire fluctuations of the amplitude in the early universe (in particular, during the inflation).  Then, the scalar field eventually decays and reheats the universe.  Importantly, in this class of scenario, the CMB radiation we observe today originates to $\\phi$ rather than to the inflaton.  As a result, one can expect that the fluctuation of the scalar-field amplitude affects the cosmic density perturbations \\cite{PRD42-313,NPB626-395,PLB524-5,PLB522-215}.  In particular, in such a scenario, adiabatic and isocurvature perturbations may be generated with cross-correlation.\\footnote {For other mechanisms of generating correlated mixture of the adiabatic and isocurvature perturbations, see \\cite{adi-iso}.} Indeed, in Ref.\\ \\cite{PLB522-215}, the CMB angular power spectrum is calculated in such a framework, and it was shown that effects of the correlated isocurvature perturbation may be large enough to be seen in the on-going and future experiments.  In addition, it is possible that the dominant part of the cosmic density perturbations observed today may originate to the primordial perturbation in the amplitude of $\\phi$ rather than to the fluctuation of the inflaton amplitude. (This kind of mechanism is sometimes called ``curvaton mechanism.'')  In this paper, we consider effects of such scalar-field condensations to the cosmic density perturbations without relying on any particular scenarios.  Assuming that the universe was once dominated by the scalar-field condensation and that the reheating occurred at a later stage with generating large amount of entropy, we study the evolutions of the cosmic density perturbations and consider the CMB angular power spectrum.  As we will see, the CMB power spectrum can be affected in various ways depending on scenarios.  Interestingly, if the decay product of the $\\phi$ field does not generate the baryon asymmetry or the CDM component, then the adiabatic and isocurvature perturbations can be both generated with cross-correlation.  In this case, the CMB power spectrum may significantly change its behavior from the adiabatic result.  In particular, the simple ``curvaton'' scenario, in which the cosmic density perturbations are generated only from the primordial fluctuation of the amplitude of $\\phi$, is severely constrained by the current observations of the CMB angular power spectrum {\\sl if the adiabatic and isocurvature perturbations are generated with cross-correlation}.  The organization of the rest of this paper is as follows.  In Section \\ref{sec:scenario}, we introduce the scenario we consider and follow the thermal history of the universe.  Evolutions of the cosmic density perturbations in such a scenario is discussed in Section \\ref{sec:evolution}.  In particular, it is discussed how the density perturbation in the radiation is affected by the primordial fluctuation of the amplitude of $\\phi$.  Then, in Section \\ref{sec:cl}, we overview the behavior of the CMB anisotropy in our scenario.  Detailed discussions on the CMB power spectrum for the cases with and without the isocurvature perturbations are given in Sections \\ref{sec:noentropy} and \\ref{sec:entropy}, respectively. Section \\ref{sec:conclusion} is devoted for the conclusion. ",
        "conclusions": "\\label{sec:conclusion} \\setcounter{equation}{0}  In this paper, we discussed the effects of the late-time entropy production due to the decay of the scalar-field condensations on the cosmic density perturbations.  If the universe is reheated by the decay of the scalar field $\\phi$, many of the components in the present universe are generated from the decay products of the $\\phi$ field.  In such a case, cosmic density perturbations are affected by the fluctuation of the amplitude of $\\phi$ which may be generated during the inflation.  If all the components in the universe originate to the decay product of $\\phi$, density perturbations generated from the primordial fluctuation of $\\phi$ becomes adiabatic.  In this case, the CMB angular power spectrum from the fluctuation of $\\phi$ becomes the usual adiabatic ones with (almost) scale-invariant spectrum.  If this becomes the dominant part of the cosmic density perturbations, then we have seen that the constraints on inflation models from observations of the CMB angular power spectrum are drastically relaxed.  If the baryon or the CDM is not generated from $\\phi$ but from a new scalar field $\\psi$, on the contrary, correlated mixture of the adiabatic and isocurvature perturbations may arise.  In particular, if the $\\psi$ field starts to oscillate much before the $\\phi$ field dominates the universe, the metric and entropy perturbations obey the model-independent relation.  In this case, the CMB angular power spectrum may be significantly affected and the shape of the resultant power spectrum depends on which component has the correlated isocurvature perturbation.  If the baryonic component has the correlated isocurvature perturbation, the density perturbations after the decay of $\\phi$ become those given in Eq.\\ (\\ref{Sb}).  In this case, $C_l$ at high multipole is more enhanced relative to that at low multipole.  On the contrary, if the CDM is not generated from $\\phi$, correlated isocurvature perturbation can be induced in the CDM component as given in Eq.\\ (\\ref{Sc}).  In this case, heights of the acoustic peaks become lower relative to the SW tail.  If there exists an contamination of these component into the total angular power spectrum, $C_l$ differs from the adiabatic one.  The important point is that, in both cases, too much correlated isocurvature perturbations become inconsistent with the present observation of the CMB power spectrum.  Even with the correlated isocurvature perturbation in the baryonic sector, heights of the acoustic peaks can be reduced by introducing uncorrelated isocurvature perturbation.  This may happen, for example, if we take account of the primordial fluctuation of the scalar field $\\psi$.  (For example, this scalar field may be the Affleck-Dine field.)  In particular, if the sizes of the correlated and uncorrelated entropy perturbations are properly chosen, the resultant CMB power spectrum becomes consistent with the present observations without the effect of the inflaton fluctuation.  Even in this case, the angular power spectrum is not exactly the same as the conventional adiabatic one and the deviation from the adiabatic result may be observed at on-going and future experiments.  In summary, if there exists correlated mixture of the adiabatic and isocurvature perturbations, the total CMB angular power spectrum can be distorted and it can be a signal of the late-time entropy production due to the decay of the scalar field condensation.  The on-going and future experiments may observe such a signal.  {\\sl Acknowledgment:} We acknowledge the use of CMBFAST \\cite{cmbfast} and RADPACK \\cite{radpack} packages for our numerical calculations. This work is supported by the Grant-in-Aid for Scientific Research from the Ministry of Education, Science, Sports, and Culture of Japan, No.\\ 12047201 and No.\\ 13740138."
    },
    "0206/astro-ph0206380_arXiv.txt": {
        "abstract": "{We present a spectral analysis of 35 GRBs detected with the HETE-2 gamma-ray detectors (the FREGATE instrument) in the energy range 7-400 keV. The GRB sample analyzed is made of GRBs localized with the Wide Field X-ray Monitor onboard HETE-2 or with the GRB Interplanetary Network. We derive the spectral parameters of the time-integrated spectra, and present the distribution of the low-energy photon index, alpha, and of the peak energy, \\ep . We then discuss the existence and nature of the recently discovered X-Ray Flashes and their relationship with classical GRBs. ",
        "introduction": "\\label{introduction}   The radiation mechanisms at work during the prompt phase of GRBs remain poorly understood, despite the observation of hundreds of GRB spectra and extensive theoretical work (e.g. \\cite{cohe97,daig98,lloy00,mesz00,pana00,pira00,zhan02}). One of the reasons for this situation is the lack of broad-band coverage of this brief phase of GRB emission (contrary to the afterglows which can be observed from hours to days after the burst). Recently, however, several instruments have extended the spectral coverage of the prompt GRB emission to the X-ray range, and to optical wavelengths in the case of GRB990123 (\\cite{aker99}), raising hopes for a better understanding of this crucial phase of GRB emission.  We present here the broad-band spectra of 35 GRBs observed by HETE-2/FREGATE in the energy range 7-400 keV. We analyse the time-integrated spectra in order to derive the distribution of their peak energies and of their low-energy spectral indices. We also discuss the existence of a possible new class of soft bursts, called X-ray flashes.  HETE-2's unique instrument suite provides broadband energy coverage of the prompt emission extending into the X-ray range. The three instruments include a gamma-ray spectrometer sensitive in the range 7-400 keV (FREGATE, \\cite{atte02}), a Wide Field X-ray Monitor sensitive in the range 2-25 keV (WXM, \\cite{kawa02}) and a CCD based Soft X-ray Camera working in the range 1-14 keV (SXC, \\cite{vill02}). In this paper we restrict our analysis to FREGATE data because this instrument, with its larger field of view, detects about two times more GRBs than WXM (the Half Width at Zero Maximum is 70$^\\circ$ for FREGATE compared to 40$^\\circ$ for the WXM) and because in most cases FREGATE data are sufficient to determine the GRB spectral parameters. We finally note that FREGATE offers for the first time a continuous coverage from 7 keV to 400 hundred keV {\\it with a single instrument}. This eliminates any possible problems caused by normalizing the responses of different instruments to one another. This characteristic of FREGATE appears essential when we try to understand whether events seen at low energies are of the same nature as classical GRBs seen at higher energies.  This work follows many previous studies which contributed to our understanding of the GRB spectral properties at gamma-ray energies (\\cite{band93,pree98}) and in hard X-rays (\\cite{stro98,fron00a,kipp01}). This paper is the first of a series devoted to the spectral analysis of the GRBs detected with HETE-2. Forthcoming papers will discuss the spectral {\\it evolution} of bright GRBs and the broad-band spectral distribution from 2 keV to 400 keV by combining the data from FREGATE and the WXM for the events which are detected by both instruments. ",
        "conclusions": "\\label{conclusion}  Our observations have two interesting consequences: they confirm that the \\ep\\ distribution is broader than previously thought (\\cite{mall95, pree00, brai00}) and they show that we do not see yet the faint end of the GRB distribution. If we assume that the correlation found by Amati et al. (2002) extends down to \\ep\\ as low as 20 keV, it would imply that the isotropic-equivalent energy radiated by a GRB with \\ep\\ = 20 keV is about 80 times smaller than the isotropic-equivalent energy radiated by a \"typical\" GRB with \\ep\\ = 200 keV. If the conclusion of Frail et al. (2000) that the total energy of GRBs is roughly constant, it implies that the jet opening angle of X-ray rich GRBs are substantially larger than the jet opening angle of \"typical\" GRBs.  Future work with HETE-2 will bring several advances in this field and should contribute to our understanding of the population of soft/faint GRBs. The continuously growing GRB sample of FREGATE should provide better statistical evidence for the effects discussed in this paper and additional clues about the possible differences between bright and faint GRBs and about the nature of X-ray rich GRBs. Joint spectral analysis with the WXM will allow more precise determinations of $\\alpha$ and \\eo\\ for X-ray rich GRBs. Finally, measuring the redshifts of a greater number of GRBs detected by HETE-2 will allow us to test the extent of the correlation between the spectral hardness of GRBs and their radiated energy in gamma-rays. X-ray rich GRBs also present an interesting challenge for future GRB missions and for observers on the ground. Future GRB missions will have to detect events which are much softer and fainter than the typical GRB population sampled by BATSE. Observers on the ground are faced with events which have fainter afterglows than the classical gamma-ray bursts.  To conclude we note that the joint detection of GRB010213 (with \\ep = 2.5 keV, \\cite{kawa03,saka03}) by WXM and FREGATE and of GRB020903 (with \\ep = 0.9 keV, \\cite{kawa03}) by WXM only demonstrate the existence of events which are even softer than the bulk of the X-ray rich GRBs discussed in this paper."
    },
    "0206/astro-ph0206455_arXiv.txt": {
        "abstract": "{We measure the average gravitational shear profile of 6 massive clusters ($M_{\\rm vir} \\sim 10^{15} M_{\\sun}$) at $z=0.3$ out to a radius $\\sim 2h^{-1}$~Mpc. The measurements are fitted to a generalized NFW-like halo model $\\rho (r)$ with an arbitrary $r \\rightarrow 0$ slope $\\alpha$. The data are well fitted by such a model with a central cusp with $\\alpha \\sim 0.9 - 1.6$ (68\\% confidence interval). For the standard-NFW case $\\alpha = 1.0$, we find a concentration parameter $c_{\\rm vir}$ that is consistent with recent predictions from high-resolution CDM N-body simulations. Our data are also well fitted by an isothermal sphere model with a softened core. For this model, our $1\\sigma$ upper limit for the core radius corresponds to a limit $\\sigma_{\\star} \\leq 0.1 {\\rm cm}^2 {\\rm g}^{-1}$ on the elastic collision cross-section in a self-interacting dark matter model. } ",
        "introduction": "Although the cold dark matter (CDM) model of structure formation has been successful in explaining many observable properties of the universe, there remain some notable discrepancies with observations. In particular, numerical simulations of dark matter haloes in CDM universes appear to overproduce the abundance of dark matter sub-haloes corresponding to dwarf galaxies by at least an order of magnitude (Klypin et al.\\ 1999; Moore et al.\\ 1999). Also, simulations indicate a transfer of angular momentum from the baryonic matter component to the non-baryonic dark matter during the assembly of disk galaxies, which leads to a conflict with the observed angular momentum properties of disk galaxies (e.g., Navarro \\& Steinmetz 1997,2000; Sommer-Larsen, G{\\\" o}tz, \\& Portinari 2002).  Finally, the simulations predict that dark matter halos in a wide mass range from dwarf galaxies to massive clusters follow an universal density profile with a central cusp $\\rho(r) \\propto r^{-\\alpha}$ with $\\alpha$ in the range $1.0 < \\alpha < 1.5$ (Navarro, Frenk, \\& White 1996, 1997; hereafter collectively NFW; Moore et al.\\ 1999). This appears to be contradicted by some observations that indicate an almost constant-density core rather than a cusp, both on galaxy scales (e.g., Dalcanton \\& Bernstein 2000; Salucci \\& Burkert 2000; de Blok et al.\\ 2001) and on cluster scales (Tyson, Kochanski, \\& Dell'Antonio 1998), although the interpretation of these results is controversial (e.g., Broadhurst et al.\\ 2000; Shapiro \\& Iliev 2000; Czoske et al.\\ 2002). This has led to suggestions that the dark matter properties may deviate from standard CDM, e.g., the dark matter could be warm (Col{\\'i}n, Avila-Reese, \\& Valenzuela 2000; Sommer-Larsen \\& Dolgov 2001), repulsive (Goodman 2000), fluid (Peebles 2000), fuzzy (Hu, Barkana, \\& Gruzinov 2000), decaying (Cen 2001), annihilating (Kaplinghat, Knox, \\& Turner 2000), self-interacting (Spergel \\& Steinhardt 2000; Yoshida et al.\\ 2000, Dav{\\'e} et al.\\ 2001) or both warm and self-interacting (Hannestad \\& Scherrer 2000). Alternatively, it has been suggested that stellar feedback from the first generation of stars formed in galaxies was so efficient that the remaining gas was expelled on a timescale comparable to, or less than, the local dynamical timescale. The dark matter subsequently adjusted to form an approximately constant density core (e.g., Gelato \\& Sommer-Larsen 1999). This is however unlikely to affect cluster cusps.  It is dangerous to draw wide-ranging conclusions about halo density profiles from a single object, even when based on high-quality data from strong gravitational lensing. Clearly, a larger number of clusters should be investigated using a range of methods. Other studies have sought to put constraints on $\\alpha$ from the observed abundances of strongly lensed arcs produced by clusters (Meneghetti et al.\\ 2001; Molikawa \\& Hattori 2001; Oguri, Taruya, \\& Suto 2001), and find best fit values of $\\alpha$ similar to CDM predictions.  Previous weak lensing measurements of cluster mass profiles have usually only considered two models, a singular isothermal sphere (SIS) model with $\\rho \\propto r^{-2}$ and the NFW model. Typically, it is found that the NFW model provides a marginally better fit to the data, but usually not at a very high significance (Clowe et al.\\ 2000; Clowe \\& Schneider 2001; Hoekstra et al.\\ 2002). Many of the clusters studied in this way show evidence for substructure, indicating that they may not be in dynamical equilibrium.  Here, we consider a density profile which is a generalized version of the NFW model, as described in \\S 2. The data we present here, detailed in \\S 3, consists of measurements of gravitational shear over a wide range of radii around 6 clusters at $z \\simeq 0.30$. In \\S 4 we present the results of our model fits, and in \\S 5 we discuss the implications for cluster-scale dark matter halos and dark matter physics. We consider both a spatially-flat $\\Lambda$CDM model with $(\\Omega_{0},\\lambda_{0}) = (0.3,0.7)$ and a SCDM (Einstein-de Sitter) cosmology with $(\\Omega_{0},\\lambda_{0}) = (1,0)$, and we use $H_0 = 100 h\\, {\\rm km}\\, {\\rm s}^{-1} {\\rm Mpc}^{-1}$. ",
        "conclusions": "We find that the average density profile of 6 massive clusters is well fitted by a generalized-NFW model with central slope $\\alpha \\sim 0.9-1.6$. This (steep) slope fits well with predictions from collisionless CDM simulations. The best fit slope of $\\alpha = 1.3-1.5$ is slightly steeper than the canonical NFW value, although $\\alpha=1$ is only disfavored at less than 90\\% confidence for any of the models. For $\\Lambda$CDM, $\\alpha \\lesssim 0.5$ can be excluded with 95\\% confidence. The inclusion of an additional SIS mass profile corresponding to a massive central galaxy only resulted in minor changes for the best fit parameters. This indicates that any additional mass component associated with the central galaxy does not make a significant contribution to the total projected density on the scales that we probe in our weak lensing study.  Our constraints on $\\alpha$ are compatible with recent strong lensing studies of massive clusters: Smith et al.\\ (2001) estimate  $\\alpha = 1.3$ for \\objectname{A383} and Gavazzi et al.\\ (2002) find $0.7 \\leq \\alpha \\leq 1.2$ for \\objectname{MS2137-23} (see, however, Sand, Treu \\& Ellis 2002, who find a best fit $\\alpha = 0.35$ and $\\alpha < 0.9$ at 99\\% CL for the same cluster). Our results are also consistent with recent X-ray cluster studies based on Chandra data: Arabadjis, Bautz, \\& Garmire (2002) find that the mass profile of \\objectname{CL 1358+6245} is well fit by a standard-NFW $\\alpha = 1$ profile, and  Lewis, Buote, \\& Stocke (2003) find $\\alpha = 1.19 \\pm 0.04$ for the central slope of \\objectname{A2029}.  We find that an isothermal sphere model with a softened core can also produce a good fit to our data. This indicates that our weak lensing data is unable to discriminate between an outer slope of $r^{-2}$ and $r^{-3}$. A pure SIS model is however strongly excluded at the 99.9\\% CL (see Fig.~\\ref{fig:NSIS}). Yoshida et al.\\ (2000) use simulations to study the effect of self-interacting dark matter on the structure of cluster halos. They find that even a small elastic collision cross-section of $\\sigma_{\\star} = 0.1 {\\rm cm}^2 {\\rm g}^{-1}$ will significantly affect the central density profile. Meneghetti et al.\\ (2001) study the strong lensing properties of the cluster models of Yoshida et al.\\ (2000), and find that even a cross-section as small as $0.1 {\\rm cm}^2 {\\rm g}^{-1}$ is incompatible with the observed abundances of strongly lensed radial and tangential arcs in clusters. Their study is based on a simulated cluster at $z = 0.278$ of final virial mass $7.4 \\times 10^{14} h^{-1} M_{\\sun}$ in a $\\Lambda$CDM universe. This cluster is at a similar redshift as the clusters in our data set, but our clusters are on average about twice as massive. From the $r_c$ values tabulated by Meneghetti et al.\\ (2001), we find that the predicted core radius for $\\sigma_{\\star} = 0.1 {\\rm cm}^2 {\\rm g}^{-1}$ of $r_c = 80h^{-1}$~kpc is similar to our $1\\sigma$ (and $2\\sigma$) upper limit for the core radius shown in Figure~\\ref{fig:NSIS}. Here, we have taken into account a factor $\\sqrt{3}$ difference in the definition of $r_c$ and a $M^{1/3}$ scaling of lengths. A cross-section of $1 {\\rm cm}^2 {\\rm g}^{-1}$ would produce a core radius of $180h^{-1}$~kpc, which is strongly excluded by our data. In the small mean free path limit of ``fluid'' dark matter, a core-collapse producing a SIS-type profile is expected. This is strongly excluded by our data, and the fluid limit is also excluded by the observed ellipticities of clusters (Miralda-Escud{\\' e} 2002).  Dav{\\'e} et al.\\ (2001) show that cross-sections of order $5 {\\rm cm}^2 {\\rm g}^{-1}$ are needed to produce a good fit to the large apparent cores in dwarf galaxies. By introducing a velocity-dependent interaction cross-section for the dark matter particles, it may still be possible to reconcile our results with a self-interacting dark matter model that can produce constant-density cores in dwarf galaxies. Even if such a model may still be allowed, its velocity-dependence is already strongly confined by other astrophysical constraints, including the demographics of supermassive black holes in galaxy nuclei (Hennawi \\& Ostriker 2001).  It should be noted that even though dark matter models which produce large constant-density cores on clusters scales appear to be ruled out, many of the non-standard dark matter models are easily compatible with cusped cluster halos. One example is warm dark matter where the thermal velocity of the dark matter particles can erase the cusps of dwarf galaxy halos, but almost not affect more massive halos at all.   \\ifthenelse{\\equal{\\version}{_working}} {"
    },
    "0206/astro-ph0206039_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{chap:intro} ",
        "conclusions": "\\label{sec:discussion}  In this chapter, using the full radiation transfer function, we have numerically computed the primary cosmic microwave background bispectrum (Eq.(\\ref{eq:almspec})) and skewness (Eq.(\\ref{eq:skewness})) down to arcminutes angular scales. As the primary bispectrum oscillates around zero (figure~\\ref{fig:bispectrum}), the primary skewness saturates at the {\\it MAP} angular resolution scale, $l\\sim 500$ (figure~\\ref{fig:skewness}). We have introduced the {\\it reduced} bispectrum, $b_{l_1l_2l_3}$, defined by equation (\\ref{eq:func}), and found that this quantity is more useful to describe physical properties of the bispectrum than the angular averaged bispectrum, $B_{l_1l_2l_3}$ (Eq.(\\ref{eq:blll})).   Figure~\\ref{fig:sn} compares the expected signal-to-noise ratio of detecting the primary non-Gaussianity based on the bispectrum (Eq.(\\ref{eq:sn})) with that based on the skewness (Eq.(\\ref{eq:skew_sn})). It shows that the bispectrum is almost an order of magnitude more sensitive to the non-Gaussianity than the skewness. We conclude that when we can compute the predicted form of the bispectrum, it becomes a ``matched filter'' for detecting the non-Gaussianity in data, and thus much more powerful tool than the skewness. Table~\\ref{tab:fnl} summarizes the minimum $f_{\\rm NL}$ for detecting the primary non-Gaussianity using the bispectrum or the skewness for {\\it COBE}, {\\it MAP}, {\\it Planck}, and the ideal experiments. This shows that even the ideal experiment needs $f_{\\rm NL}>3$ to detect the primary bispectrum.  \\begin{figure} \\plotone{figure5.eps} \\caption{Bispectrum vs Skewness} \\mycaption{Comparison of the signal-to-noise ratio summed up to a certain $l_3$, $S/N(<l_3)$, for the bispectrum (top panel; Eq.(\\ref{eq:sn})) and the skewness (bottom panel; Eq.(\\ref{eq:skew_sn})). $S/N(<l_3)$ is in units of $f_{\\rm NL}$. The dotted lines show {\\it COBE}, {\\it MAP}, and {\\it Planck} experiments (dotted lines), while the solid line shows the ideal experiment. See table~\\ref{tab:fnl} for $f_{\\rm NL}$ to obtain $S/N>1$.} \\label{fig:sn} \\end{figure}  \\begin{table} \\caption{Detection Limit for the Non-linear Coupling Parameter} \\mycaption{The minimum non-linear coupling parameter, $f_{\\rm NL}$, needed for detecting the primary non-Gaussianity by the bispectrum or the skewness with the signal-to-noise ratio greater than 1. These estimates include the effects of cosmic variance, detector noise, and foreground sources.} \\label{tab:fnl} \\begin{center} \\begin{tabular}{ccc}\\hline\\hline Experiments & $f_{\\rm NL}$ (Bispectrum) & $f_{\\rm NL}$ (Skewness) \\\\ \\hline {\\it COBE}   & 600 & 800 \\\\ {\\it MAP}    & 20  & 80 \\\\ {\\it Planck} & 5   & 70 \\\\ Ideal  & 3   & 60 \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}  We have calculated the secondary bispectra from the coupling between the SZ effect and the weak lensing effect, and from the extragalactic radio and infrared sources. Only {\\it Planck} will detect the SZ--lensing bispectrum, while both {\\it MAP} and {\\it Planck} will detect the extragalactic point-source bispectrum (table~\\ref{tab:sn}).   We have also studied how well we can discriminate between the primary, the SZ--lensing coupling, and the extragalactic point-source bispectra. We have found that {\\it MAP} and {\\it Planck} will separate the primary from the other secondary sources with 1\\% or better accuracy. This conclusion is due to the acoustic oscillation in the primary bispectrum that does not appear in the secondary bispectra. The SZ--lensing coupling and the extragalactic sources are well separately measured by {\\it Planck} experiment, although {\\it COBE} and {\\it MAP} cannot discriminate between them (table~\\ref{tab:corr}).   Our arguments on the ability to discriminate between various bispectra have been based upon the shape difference, and thus have not taken into account the spectral difference in frequency space. \\citet{TE96} and \\citet{CHT00} have shown that the multi-band observation is efficient for discriminating between the primary signal and the other foreground sources in the CMB power spectrum. Their scheme should be effective on the bispectrum as well, and will improve the accuracy of the foreground removal further. We thus expect that {\\it MAP} and {\\it Planck} will measure the primary bispectrum separately from the foregrounds.   Simple slowly-rolling single-field inflation models predict $f_{\\rm NL}\\sim {\\cal O}(10^{-2})$ \\citep{SB90,SB91,Gan94}, while the second order perturbation theory predicts $f_{\\rm NL}\\sim {\\cal O}(1)$ \\citep{PC96}; thus, significant detection of the primary bispectrum or the skewness with any experiments means that these inflation models need to be modified. According to our results, if the reported detection \\citep{FMG98,Mag00} of the bispectrum on the {\\it COBE} map were cosmological in origin, then {\\it MAP} and {\\it Planck} would detect the primary bispectrum much more significantly. While Banday, Zaroubi and G\\'orski \\citep{BZG00} have shown that the one of those detections \\citep{FMG98} is accounted for by the experimental systematic effects, the other \\citep{Mag00} is significant even after removing such the systematics.   Although we have not discussed so far, spatial distribution of interstellar dust emissions is a potential source of microwave non-Gaussianity. While it is very hard to estimate the bispectrum analytically, we can use the dust template map compiled by \\citet{SFD98} to estimate the dust bispectrum. For example, we have found that the dimensionless skewness parameter, $\\left<(\\Delta T)^3\\right>/\\left<(\\Delta T)^2\\right>^{3/2}$, is as large as 51 on the template map. We have used the publicly available HEALPix-formatted \\citep{GHW98} $100~\\mu{\\rm m}$ map, which contains 12,582,912 pixels without sky-cut. The mean intensity in the map is $14.8~{\\rm MJy~sr^{-1}}$. Of course, this skewness is largely an overestimate for the real CMB measurement; we need to cut a fraction of the sky that contains the Galactic plane, and this will greatly reduce the non-Gaussianity. Nevertheless, residual non-Gaussianity is still a source of the microwave bispectrum, and has to be taken into account. Moreover, the form of the measured bispectrum on the dust map reflects the physics of interstellar dust, which is highly uncertain at present; thus, studying the interstellar dust bispectrum is a challenging field.   \\chapter{Measurement of Bispectrum on the {\\it COBE} DMR sky maps} \\label{chap:obs_bl}   Several authors have attempted to measure non-Gaussianity in CMB using various statistical techniques (e.g., Kogut et al. 1996b); as yet no conclusive detection has been reported except for measurement of several modes of the normalized CMB bispectrum on the {\\it COBE} Differential Microwave Radiometer (DMR) sky maps \\citep{FMG98,Mag00}. The existence of non-Gaussianity in the DMR data is controversial. If the CMB sky were non-Gaussian, this would challenge our simplest inflationary model.   The angular bispectrum, $B_{l_1l_2l_3}$, is the harmonic transform of the three-point correlation function. We carefully distinguish the normalized bispectrum, $B_{l_1l_2l_3}/\\left(C_{l_1}C_{l_2}C_{l_3}\\right)^{1/2}$, from the bispectrum, $B_{l_1l_2l_3}$. \\citet{FMG98} have measured 9 equilateral ($l_1=l_2=l_3$) modes of the normalized bispectrum, $B_{l_1l_2l_3}/\\left(C_{l_1}C_{l_2}C_{l_3}\\right)^{1/2}$, on the DMR map, claiming detection at $l_1=l_2=l_3=16$. Their result has been under extensive efforts to confirm its significance and origin. \\citet{BT99} claim that a few individual pixels in the DMR map are responsible for the most of the signal. \\citet{BZG00} have proposed an eclipse effect by the Earth against the {\\it COBE} satellite as a possible source of the signal. \\citet{Mag00} has measured other 8 inter-$l$ modes of the normalized bispectrum such as $B_{l-1ll+1}/\\left(C_{l-1}C_{l}C_{l+1}\\right)^{1/2}$, and claims that scatter of the normalized bispectrum among 8 modes is too small to be consistent with Gaussian. \\citet{SM00} further report measurement of 24 other inter-$l$ modes for different lags in $l$, and conclude they are consistent with Gaussian.   Thus, until now 41 modes of the normalized bispectrum have been measured on the DMR map. Here, we simply ask: ``how many modes are available in the DMR map for the bispectrum?'' The answer is 466, up to a maximum multipole of 20 that corresponds to the DMR beam size; thus, it is conceivable that the claimed detection of the normalized bispectrum at $l_1=l_2=l_3=16$ would be explained by a statistical fluctuation, as 9 modes are expected to have statistical significance above 98\\% out of 466 independent modes even if CMB is exactly Gaussian. In this chapter, we measure 466 modes of the CMB bispectrum on the {\\it COBE} DMR sky maps, testing the claimed detection of the bispectrum and non-Gaussianity. We take into account the covariance between these modes due to the Galactic cut, which has not been done in the previous work.   On the theoretical side, several predictions for the CMB bispectrum exist. Several authors \\citep{FRS93,LS93,Gan94} have predicted the primary bispectrum (or equivalently three-point correlation function) on the DMR angular scales from slow-roll inflation models. In chapter~\\ref{chap:theory_bl}, we have extended the prediction down to arcminutes scales using the full radiation transfer function.   In addition to the primary one, secondary sources in the low-redshift universe and foreground sources produce the bispectrum through their non-linearity. \\citet{LS93} and \\citet{SG99} have calculated the secondary bispectrum arising from non-linear evolution of gravitational potential; \\citet{GS99} and \\citet{CH00} have calculated the one from the gravitational lensing effect coupled with various secondary anisotropy sources. In chapter~\\ref{chap:theory_bl}, we have calculated the foreground bispectrum from extragalactic radio and infrared point sources. While the bispectrum is not the best tool for detecting the signature of rare highly non-linear events, e.g., textures \\citep{PK01}, it is sensitive to weakly non-linear effects.   Having theoretical predictions is a great advantage in extracting physical information from measurement; one can fit a predicted bispectrum to the data so as to constrain parameters in a theory. Since the DMR beam size is large enough to minimize contribution from the secondary and the extragalactic foreground sources, the only relevant source would be the primary one. In this chapter, we fit a theoretical primary bispectrum \\citep{KS01a} to the data.   The Galactic plane contains strong microwave emissions from interstellar sources. The emissions are highly non-Gaussian, and distributed on fairly large angular scales. Unfortunately, predicting the CMB bispectrum from interstellar sources is very difficult; thus, we excise the galactic plane from the DMR data. We model the residual foreground bispectrum at high galactic latitude using foreground template maps. By simultaneously fitting the foreground bispectrum and the primary bispectrum to the DMR data for three different Galactic cuts, we quantify the importance of the interstellar emissions in our analysis.   This chapter is organized as follows. In \\S~\\ref{sec:bispectrum*}, we define the angular bispectrum, and show how to compute it efficiently from observational data. In \\S~\\ref{sec:measurement}, we study statistical properties of the bispectrum and the normalized bispectrum. We then measure the normalized bispectrum from the {\\it COBE} DMR four-year sky maps \\citep{Ben96}, testing Gaussianity of the DMR map. In \\S~\\ref{sec:fit}, we fit predicted bispectra to the DMR data, constraining parameters in the predictions. The predictions include the primary bispectrum from inflation and the foreground bispectrum from interstellar Galactic emissions. Finally, \\S~\\ref{sec:discussion_bl} concludes.   \\label{sec:discussion_bl}  In this chapter, we have measured all independent configurations of the angular bispectrum on the {\\it COBE} DMR map, down to the DMR beam size. Using the most sensitive sky map to CMB, which combines the maps at 53 and 90~GHz, we test the Gaussianity of the DMR map.   We find that the normalized bispectrum, $B_{l_1l_2l_3}/\\left(C_{l_1}C_{l_2}C_{l_3}\\right)^{1/2}$, gives more robust test of Gaussianity than the bare bispectrum, $B_{l_1l_2l_3}$. We compare the measured data with the simulated realizations, finding the DMR map comfortably consistent with Gaussian. We explain the reported detection of the normalized bispectrum at $l_1=l_2=l_3=16$ \\citep{FMG98} by a statistical fluctuation as an alternative to the \"eclipse effect\" proposition made in \\citet{BZG00}.   We fit the predicted bispectra to the data, constraining the parameters in the predictions, which include the primary bispectrum from inflation and the foreground bispectra from interstellar dust and synchrotron emissions. We find that neither dust nor synchrotron emissions contribute to the bispectrum significantly.   We have obtained a weak constraint on the non-linear coupling parameter, $f_{\\rm NL}$, that characterizes non-linearity in inflation. We interpret the constraint in terms of a single-field inflation as follows. According to the analysis of non-linear perturbations on super horizon scales \\citep{SB90}, we can explicitly calculate $f_{\\rm NL}$ as \\begin{equation} \\label{eq:SB91} f_{\\rm NL}= -\\frac{5}{24\\pi G} \\left(\\frac{\\partial^2\\ln H}{\\partial\\phi^2}\\right), \\end{equation} where $H$ is the Hubble parameter during inflation. When applying the slow-roll conditions to an inflaton potential $V(\\phi)$, we have $\\partial\\ln H/\\partial\\phi\\approx (d\\ln V/d\\phi)/2$; thus, $f_{\\rm NL}$ is on the order of curvature of a slow-roll potential, implying that $\\left|f_{\\rm NL}\\right|$ should be smaller than 1 in slow-roll inflation. Therefore, the obtained constraint, $\\left|f_{\\rm NL}\\right|<1.6\\times 10^3$, seems too weak to be interesting; however, any deviation from slow-roll could yield larger $\\left|f_{\\rm NL}\\right|$, bigger non-Gaussianity.   The next generation satellite experiments, {\\it MAP} and {\\it Planck}, should be able to put more stringent constraints on $f_{\\rm NL}$. In chapter~\\ref{chap:theory_bl}, we have shown that {\\it MAP} and {\\it Planck} should be sensitive down to $\\left|f_{\\rm NL}\\right|\\sim 20$ and 5, respectively. We find that the actual constraint from {\\it COBE} (figure~\\ref{fig:fNL}) is much worse than the estimate. This is partly due to different cosmology used for the model, but mainly due to incomplete sky coverage; the statistical power of the bispectrum at low multipoles is significantly weakened by the Galactic cut. Since {\\it MAP} and {\\it Planck} probe much smaller angular scales, and their better angular resolution makes an extent of the Galactic cut smaller, the degradation of sensitivity should be minimal. Moreover, the improved frequency coverage of future experiments will aid in extracting more usable CMB pixels from the data. At this level of sensitivity, any deviation from slow-roll could give an interesting amount of the bispectrum, and {\\it MAP} and {\\it Planck} will put severe constraints on any substantial deviation from slow-roll.   While we have explored adiabatic generation of the bispectrum only, isocurvature perturbations from inflation also generate non-Gaussianity \\citep{LM97,P97,BZ97}. They are in general more non-Gaussian than the adiabatic perturbations; it is worth constraining these models by the same strategy as we have done in this chapter.   \\chapter{In Pursuit of Angular Trispectrum} \\label{chap:obs_tl}   The angular trispectrum, the harmonic transform of the angular four-point correlation function, carries cosmological information which is independent of the power spectrum and the bispectrum.   Several authors show that the weak lensing effect produces non-trivial connected CMB angular trispectrum or four-point correlation function on small angular scales \\citep{Ber97,ZS99,Zal00,Hu01}. On large angular scales, \\citet{Ino01a,Ino01b} shows that the connected trispectrum is produced if topology of our universe is closed flat or closed hyperbolic. These effects do not produce the bispectrum, the angular three-point harmonic spectrum, but the trispectrum. Hence, while we have not found significant bispectrum on the DMR map in chapter~\\ref{chap:obs_bl}, we could find the trispectrum.   The connected angular trispectrum contributes to the power spectrum covariance matrix. It increases the variance, and produces non-zero off-diagonal terms in the covariance matrix. Measuring the connected trispectrum thus constrains how much the connected trispectrum contributes to the power spectrum covariance. This is important to do. We have to know the power spectrum covariance matrix as accurate as possible, when we measure the power spectrum with better than 1\\% accuracy, and determine many of cosmological parameters with better than 10\\% accuracy.   So far, there has been no attempt to measure the angular trispectrum of the CMB anisotropy. In this chapter, we present the first measurement of the CMB trispectrum on the DMR map. We measure all independent terms of the angular trispectrum down to the DMR beam size, and test Gaussianity of the DMR data.   This chapter is organized as follows. In \\S~\\ref{sec:trispectrum*}, we give our method of measuring the angular trispectrum from CMB sky maps. In \\S~\\ref{sec:norm_tl}, we study statistical properties of the normalized trispectrum. In \\S~\\ref{sec:test}, applying the method to the DMR data, we measure the angular trispectrum. We then test Gaussianity of the DMR data with the normalized trispectrum. Finally, \\S~\\ref{sec:conclusion_obstl} concludes this chapter,   \\label{sec:conclusion_obstl}  In this chapter, we have presented the first measurement of the CMB angular trispectrum on the {\\it COBE} DMR sky maps. We have measured all the trispectrum terms, 21,012 terms, down to the DMR beam size. Since 190 $L=0$ modes have no statistical power of testing Gaussianity, we have used 20,822 $L\\neq 0$ modes to test Gaussianity of the DMR data, and found that the DMR map is comfortably consistent with Gaussianity.   Our results do not directly constrain the connected trispectrum for $L=0$, $T^{ll}_{l'l'}(0)$, which contributes to the power spectrum covariance through equation~(\\ref{eq:covpowerspec}). We can thus conclude nothing as to whether the covariance matrix is diagonal on the DMR angular scales. Moreover, $T^{ll}_{ll}(0)$ increases the power spectrum variance; we have no idea how much the contribution is. We need to use other statistics than the angular trispectrum to investigate the power spectrum covariance. Otherwise, we have to have a model for the connected trispectrum, and constrain $T^{ll}_{l'l'}(0)$ by measuring the other trispectrum configurations.   One example for a trispectrum model is the one produced in a closed hyperbolic universe. \\citet{Ino01b} suggests that the closed hyperbolic geometry produces non-zero connected trispectrum. In appendix~\\ref{app:CH}, we have derived an analytic prediction for the connected trispectrum produced in a closed hyperbolic universe (Eq.(\\ref{eq:CHprediction})). For $L=0$, we reduce the prediction to \\begin{equation} \\left<T^{l_1l_1}_{l_3l_3}(0)\\right>_{\\rm c} = (-1)^{l_1+l_3}\\sqrt{(2l_1+1)(2l_3+1)}\\left(2F_{l_1l_3}\\right) + 4 F_{l_1l_1}\\delta_{l_1l_3}, \\end{equation} where \\begin{equation} F_{l_1l_3}= \\sum_{\\nu}P_\\Phi^2(\\nu)g_{{\\rm T}l_1}^2(\\nu)g_{{\\rm T}l_3}^2(\\nu) \\left<\\left|\\xi_{l_1m_1}(\\nu)\\right|^2\\right> \\left<\\left|\\xi_{l_3m_3}(\\nu)\\right|^2\\right>, \\end{equation} and $\\nu=\\sqrt{k^2-1}$ is a discrete wavenumber, $\\Phi(\\nu)$ is the primordial curvature perturbation, $g_{{\\rm T}l}(\\nu)$ is the radiation transfer function, and $\\xi_{lm}(\\nu)$ describes eigenmodes in a closed hyperbolic geometry. Hence, a single function, $F_{ll'}$, determines the connected trispectrum completely.   $F_{ll'}$ contributes to the power spectrum covariance directly. Using the prediction, we find the power spectrum covariance matrix in a closed hyperbolic universe (Eq.(\\ref{eq:CHcov})), \\begin{equation} \\left<C_lC_{l'}\\right> - \\left<C_l\\right>\\left<C_{l'}\\right> = \\frac{2}{2l+1}\\left(\\left<C_l\\right>^2+2F_{ll}\\right)\\delta_{ll'} +2F_{ll'}. \\end{equation} The power spectrum is given by \\begin{equation} \\left<C_l\\right>= \\sum_{\\nu}P_\\Phi(\\nu)g_{{\\rm T}l}^2(\\nu) \\left<\\left|\\xi_{lm}(\\nu)\\right|^2\\right>. \\end{equation} Note that $\\left<C_l\\right>\\left<C_{l'}\\right>\\ge F_{ll'}$. It follows from this equation that to constrain the contribution of the connected $T^{ll}_{l'l'}(0)$ to the power spectrum covariance, we need to constrain $F_{ll'}$.   In addition to $L=0$ modes, fortunately, a closed hyperbolic universe also produces the connected trispectrum for the group (c), $l_1=l_2=l_3=l_4\\equiv l$ and $L\\neq 0$. The prediction is $\\left<T^{ll}_{ll}(L)\\right>_{\\rm c}=4(2L+1)F_{ll}$. We may constrain this term with our measurement of the group (c). The group (c) is, however, the most noisy group, and the constraint from the DMR measurement is too weak to be useful yet.   The most promising way to investigate the angular trispectrum in near future is to use the {\\it MAP} CMB sky maps. The big advantage of {\\it MAP} over {\\it COBE} is the much higher angular resolution. The high-resolution measurement is important even on large angular scales, as we can minimize the effect of the Galactic cut. Since the rather big {\\it COBE} Galactic cut has been the main cause of numerous complications in the data analysis, we expect that the {\\it MAP} data will measure the angular trispectrum with much better sensitivity, and with much smaller systematic errors than {\\it COBE}. The {\\it MAP} trispectrum will put strong constraints on several challenging non-Gaussian models: the non-Gaussianity induced from topology of the universe, the weak lensing effect, and so on, which are not probed by the angular bispectrum."
    },
    "0206/astro-ph0206275_arXiv.txt": {
        "abstract": "{Following the discovery of spiral structure in IC3328 (Jerjen et al.~2000), we present further evidence that a sizable fraction of bright early-type dwarfs in the Virgo cluster are genuine disk galaxies, or are hosting a disk component. Among a sample of 23 nucleated dwarf ellipticals and dS0s observed with the Very Large Telescope in $B$ and $R$, we found another four systems exhibiting non-axisymmetric structures, such as a bar and/or spiral arms, indicative of a disk (IC0783, IC3349, NGC4431, IC3468). Particularly remarkable are the two-armed spiral pattern in IC0783 and the bar and trailing arms in NGC4431. For both galaxies the disk nature has recently been confirmed by a rotation velocity measurement (Simien \\& Prugniel 2002). Our photometric search is based on a Fourier decomposition method and a specific version of unsharp masking. Some ``early-type'' dwarfs in the Virgo cluster seem to be former late-type galaxies which were transformed to early-type morphology, e.g.~by ``harassment'', during their infall to the cluster, while maintaining part of their disk structure. ",
        "introduction": "The physical nature and origin of early-type dwarf (essentially dE) galaxies is still largely unknown (for a review see Ferguson \\& Binggeli 1994). According to their apparent flattening, which is similar to the one measured for giant ellipticals (Binggeli \\& Popescu 1995; Ryden \\& Terndrup 1994), early-type dwarfs seem to be spheroids. This is supported by the fact that most of these systems are not rotation-supported (Ferguson \\& Binggeli 1994; de Rijcke et al. 2001; Geha et al. 2001). On the other hand, it has always been suspected that a considerable number of early-type dwarfs might be disk galaxies. In fact, around 20 bright early-type dwarfs in the Virgo cluster are classified dS0 because of their S0-like morphology, exhibiting a lens, a bar, or high flattening (Sandage \\& Binggeli 1984; Binggeli \\& Cameron 1991). In certain evolutionary scenarios, like ram-pressure stripping (Gunn \\& Gott 1972; Abadi et al. 1999) or galaxy harassment (Moore et al. 1998), dwarf ellipticals are believed to have originated from late-type spirals or irregulars, hence some of them might have retained their disk nature. In addition, Ryden et al. (1999) showed that many early-type dwarfs have ``disky'' isophotes, similar to giant ellipticals.  Recently, Jerjen et al.~(2000) discovered a weak spiral structure in the seemingly normal dwarf elliptical galaxy IC3328 by means of deep VLT-photometry. This is unambiguous evidence for the presence of a disk in this galaxy, supporting the conjecture that the number of (hidden) disk galaxies among bright early-type dwarfs could be quite high.  Following the work of Jerjen et al. (2000), we carefully searched a larger sample of dEs observed with the VLT for additional indications of spiral or bar structure. We first applied the same techniques as for IC3328, i.e.~relying on residual images and Fourier analysis. Among the 23 dEs studied we found seven promising candidates for hidden disk structure. However, when we further explored these objects we realized that the observed spiral and bar features which we took as disk signatures are accompanied, and could actually be {\\em caused}, by a specific behaviour of the ellipticity and position angle profiles. By performing a set of simulations with artificial galaxies we convinced ourselves that the interplay between photometric parameters can indeed produce amazingly spiral-like twisting isophotes and thus mimic a genuine spiral structure. Hence, although the particular parameter combinations found in many galaxies are remarkable in themselves, they cannot unambiguously be interpreted as signs of disk structure.  Fortunately, the ambiguity can be solved by applying an unsharp masking technique, which is a model-free method to amplify {\\em local}\\/ image residuals. Only four of the seven candidates mentioned (two of which are already classified as dS0), in addition to IC3328 (Jerjen et al.~2000), withstood the scrutiny of unsharp masking: IC0783, IC3349, NGC4431, and IC3468. These are the galaxies focused on in the present paper. Although none of the four discovered cases is as spectacular as IC3328 with its tightly wound spiral, they show that the fraction of disk galaxies among bright early-type dwarfs, at least in the Virgo cluster, is 20\\% or larger. We recall that the situation is quite similar to classical (non-dwarf) elliptical galaxies, where also some Es turned out to be barred S0s, or to contain such a component, as betrayed by inner isophotal twists (Nieto et al.~1992). \\begin{table*}[t] \\caption[]{Basic data of the early-type dwarfs considered in this study} \\vspace{0.3 cm} \\begin{center} \\begin{tabular}{lllrrrrrr} \\hline \\multicolumn{1}{l}{VCC} & \\multicolumn{1}{l}{Name} & \\multicolumn{1}{l}{Type} & \\multicolumn{1}{c}{$R_T$} & \\multicolumn{1}{c}{$M_{R_T}$} & \\multicolumn{1}{l}{$B-R$} & \\multicolumn{1}{c}{$v_{\\odot}$} & \\multicolumn{1}{l}{$\\epsilon$} & \\multicolumn{1}{c}{$pa$} \\\\ \\multicolumn{1}{l}{(1)} & \\multicolumn{1}{l}{(2)} & \\multicolumn{1}{l}{(3)} & \\multicolumn{1}{c}{(4)} & \\multicolumn{1}{c}{(5)} & \\multicolumn{1}{c}{(6)} & \\multicolumn{1}{c}{(7)} & \\multicolumn{1}{c}{(8)} & \\multicolumn{1}{c}{(9)} \\\\ \\hline \\\\ 0490 & IC0783  & dS0(3),N  & 12.63 & -18.52 & 1.34 & 1293 & 0.25 & 130 \\\\ 0940 & IC3349  & dE1,N     & 13.56 & -17.59 & 1.25 & 1563 & 0.22 &  14 \\\\ 1010 & NGC4431 & dS0(5),N  & 12.47 & -18.68 & 1.39 &  913 & 0.38 & 168 \\\\ 1422 & IC3468  & E1,N:     & 12.64 & -18.51 & 1.16 & 1372 & 0.17 & 150 \\\\ \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  The rest of the paper is organized as follows. In Sect.~2 we give some observational background and put this investigation into context with our more general project. In Sect.~3 we employ, and present the results of, a Fourier analysis of the galaxy images. Sect.~4 contains a brief account of the application of unsharp masking, which is the decisive detection tool. In Sect.~5 we discuss our findings and what they could mean, case by case. A summary is given in Sect.~6. Throughout this paper we assume a distance to the Virgo cluster of $D = 17$ Mpc, corresponding to $(m-M) = 31.15$. ",
        "conclusions": "Following the discovery of spiral structure in a early-type dwarf galaxy classified as dE,N by Jerjen et al.~(2000), we searched our whole VLT-sample of bright Virgo dEs and dS0s for further photometric disk signatures. The principal search tools applied were (1) looking for inner residual features by subtracting a model image of the galaxy based on its mean surface brightness profile, but with fixed ellipticity and position angle, from the observed image; (2) a Fourier expansion of the galaxy image in polar form, where the lowest order of non-axisymmetry is indicating the presence of a bar or two-armed spiral structure; and (3) unsharp masking, where the original image is divided by an appropriately smoothed one to enhance any local ``irregularities''. Unsharp masking turned out to be the most reliable method to uncover hidden structures. The first method can give misleading results, if the inherent assumptions on global symmetry, as well as the high sensitivity of the outcome to slightly varying position angles are not taken into account. But once the clear-cut choice by unsharp masking is made, the first two methods are useful to visualize and quantify the symmetric residual structures found.  In addition to IC3328 (Jerjen et al.~2000), we found photometric traces of a possible disk component in four more early-type dwarfs out of a sample of 23:  {\\bf IC0783}: The two spiral arms of this ``dS0'' galaxy are already evident in the direct optical image. Obviously this is a disk galaxy, which now is also confirmed by the measurement of its rotation (Simien \\& Prugniel 2002). The central structure of this galaxy remains unresolved; we think there could be an inner ring.  {\\bf IC3349}: Fourier analysis and unsharp masking reveal a long and elongated (if only weak) structure in the central part, which we interpret as a bar in a nearly face-on disk. A revised type for this ``dE'' galaxy would be dSB0.  {\\bf NGC4431}: The quite strong bar present in this galaxy is the most striking discovery and clearly reveals the disk nature of this dwarf -- again nicely confirmed by its measured rotation (Simien \\& Prugniel 2002). Besides the bar we clearly note trailing arms and two dense regions on the leading side of the bar. This so-called T-structure is very similar to the results of a simulation presented by Patsis \\& Athanassoula (2000). A more fitting type for this ``dS0'' galaxy would be dSB0/a.  {\\bf IC3468}: In the very center of this dwarf elliptical we either observe a rather short bar in a nearly face-on disk, or a small disk seen edge-on in a spheroid. Surprisingly, Simien \\& Prugniel (2002) found essentially zero rotation along the position angle of this structure, which renders a clear interpretation of what we see impossible at present.  We emphasize that none of these objects is comparable to IC3328. The weak and uniformly wound spiral structure in this galaxy seems to be truly unique, constituting a particular class of dwarf galaxies; at least we did not find an additional example.  Our findings confirm previous suggestions that a sizeable fraction of all bright early-type dwarfs in the Virgo cluster are disk galaxies. In a possible scenario for their evolution they are former late-type disk galaxies which have been transformed to the systems we observe today during their infall to the cluster. The discovery of more objects of this kind in the Virgo cluster, but also in other clusters, could therefore further constrain possible models for the formation and evolution of early-type galaxies in general."
    },
    "0206/astro-ph0206043_arXiv.txt": {
        "abstract": "We compare in detail the results of simulations of electromagnetic showers in ice in the GeV-TeV energy range, using both the GEANT package and the ZHS Monte Carlo, a code specifically designed to calculate coherent \\v Cerenkov radio pulses from electromagnetic showers in dense media. The longitudinal and lateral profiles as well as the tracklengths, and excess tracklengths are shown to agree at the $10\\%$ level. We briefly comment on the negligible influence of the Landau-Pomeran\\v cuk-Migdal effect on the total shower tracklength. Our results are relevant for experiments exploiting the radio \\v Cerenkov technique in dense media as well as for other detectors that rely on the \\v Cerenkov effect in water or ice. ",
        "introduction": "Ultra high energy neutrino (UHE$\\nu$) detection is one of the experimental fields in astroparticle physics that has received most attention in the last two decades \\cite{gaisser95}. Several experiments are already taking data and several more are under way or in the proposal stage. Due to the low neutrino interaction probability and the low expected neutrino fluxes, immense volumes of detector material ($\\sim 1~{\\rm km}^3$ at least) are required \\cite{halzen00}. Currently explored detection techniques exploit the observation of \\v Cerenkov radiation in the optical frequency range from neutrino-induced showers and neutrino-induced charged leptons in dense, transparent media \\cite{halzen00,andres01,hooper02} or the search for horizontal air showers \\cite{cronin98}.  The observation of coherent \\v Cerenkov pulses in the MHz-GHz frequency range from neutrino induced showers in transparent, dense media, provides an alternative method of detecting UHE$\\nu$'s \\cite{radhep2000}. The technique is most promising at neutrino energies $>$PeV ($10^{15}$ eV) at which effective volumes in excess of $1~{\\rm km}^3$ can be achieved in a cost-effective manner \\cite{price96,frichter96}. Electromagnetic showers are known to develop an excess of negative charge of about $20\\%$ mainly due to Compton scattering. When the wavelength of the radiation is larger than the typical dimensions of the shower the emission from the excess charge is coherent, and the power in radio waves scales as the square of shower energy as predicted by Askary'an in the 1960's~\\cite{askaryan61}.  As interest in high energy neutrino detection grew in the mid 1980's, proposals were made to search for radio pulses produced by the showers that develop when UHE neutrinos interact. Arrays of antennas could be used to detect the pulses produced in deep ice \\cite{markov86} and radiotelescopes for those produced under the Moon surface \\cite{zheleznykh89}. Clearly a reliable calculation of radio pulses was needed to explore the possibilities of these techniques. A fast Monte Carlo code to simulate up to PeV electromagnetic showers was specifically designed in the early 1990's to calculate the interference pattern in the MHz-GHz region from an electromagnetic shower in ice (the ZHS code from now on) \\cite{zas91,zas92}. The results revealed the wealth of information that is kept in the radiation pattern \\footnote{As the radiation is coherent, it is in principle possible to extract the spatial charge distribution from the electric field amplitude.}. Perspectives for the technique became most encouraging for energies above the 10~PeV scale. The radio technique was further studied in the 1990's from three perspectives: prospects for radio detection arrangements were discussed \\cite{frichter96}, experimental measurements were performed \\cite{RICE01}, and the pulse simulations were extended to the most promising EeV range where full simulations are out of question with conventional computing facilities \\cite{alvarez97,alvarez98,alvarez99,alvarez00}.  The existence of the charge excess was recently confirmed in a revealing accelerator experiment at SLAC \\cite{saltzberg00}, giving a good thrust to the technique and generating a large number of proposals, some of them already in operation \\cite{gorham-aspen}. Two experiments are presently active: the RICE experiment, an array of antennas buried in the transparent polar ice cap \\cite{RICE01} and the GLUE experiment which uses the visible side of the Moon as target for UHE$\\nu$ and cosmic ray interactions \\cite{GLUE00}. New proposals include: ANITA~\\cite{gorham-aspen}, a balloon flown antenna looking down to the polar ice cap; SalSA~\\cite{SalSA01}, the same concept as RICE but exploiting the excellent optical properties of some salt domes; a proposal to analyze data from the FORTE satellite~\\cite{FORTE} \\footnote{FORTE searches for radio transients related to weather and has already triggered several million pulses in the 30-300 MHz frequency range}; and the LOFAR project, which will use a radio-astronomy array of low frequency antennas to search for radio emission from extensive air showers~\\cite{LOFAR}.  Radio patterns have been studied for long by an independent group, both from the theoretical \\cite{buniy02} and the phenomenological sides \\cite{frichter96,razzaque01}. Recently the predictions of ZHS have been challenged by shower simulations performed with the all purpose GEANT 3.21 package which quote as much as 40\\% discrepancies with ZHS in relevant parameters for radio-emission \\cite{razzaque01}. The main discrepancy is on the absolute normalization of the amplitude of the electric field, a quantity that is known to be directly related to the difference in tracklength between electrons and positrons in the shower \\cite{zas91,zas92}. Since the power in the signal scales with the square of the electric field amplitude, the discrepancy is unacceptable. The simulation of high energy showers in dense media is a cornerstone to interpret future measurements and a crucial tool to optimize the experiment's performance. Clearly the discrepancies are numerically very important, they need to be understood and the relative merits of alternative calculations have to be clearly addressed.  In this paper we compare the ZHS code in the energy range between 100 GeV and 10 TeV to simulations performed with both the GEANT 3.21 package and a more recent version of GEANT, namely GEANT4. We find GEANT4 predictions to agree with those of ZHS to about 10\\%. Moreover, provided that particular care is taken in defining some internal variables in the GEANT 3.21 package, we obtain results consistent with both GEANT 4 and ZHS. We understand that our result solves the long standing discrepancies between the two groups, and gives confidence on the calculations, which is essential for the future interpretation of data. Our results may be also of interest to neutrino experiments which are sensitive to tracklengths in water, such as AMANDA \\cite{AMANDA}, BAIKAL \\cite{Baikal}, ANTARES \\cite{Antares}, SNO \\cite{SNO}, SuperKamiokande \\cite{SK}, and even to cosmic ray experiments using water \\v Cerenkov tanks such as Auger \\cite{auger}.  This paper is organized as follows. Firstly we discuss the basis of radio interference calculations and remind the reader of the physics involved in the coherent \\v Cerenkov radio emission from showers. In section \\ref{sec:MC} we briefly discuss the inputs of the ZHS and GEANT simulators, and we quantitatively discuss the differences between the simulations. We also explore the main differences between the codes, namely the implementation of the LPM effect in the GeV-TeV range and its influence on shower development. Section \\ref{sec:conclusions} concludes the paper. ",
        "conclusions": "\\label{sec:conclusions}  We have simulated electromagnetic showers in ice using two different Monte Carlo codes, GEANT and ZHS. The discrepancies discussed in reference \\cite{razzaque01} between the ZHS code and GEANT diminish to the 10 \\% level when the GEANT 4 version is used, and also when the GEANT 3.21 code is carefully set to allow for a constant energy threshold.  We have shown that the influence of the LPM effect on shower development is negligible for energies below $E_{\\rm LPM}$, despite the fact that the bremsstrahlung cross section is significantly affected. As it has been now well established, the shower longitudinal profile is dramatically affected by the LPM at energies above $E_{\\rm LPM}$.  It is worth stressing that ZHS is a Monte Carlo specifically designed to calculate coherent radiopulses in electromagnetic showers. On the other hand, GEANT is a general purpose Monte Carlo designed to deal with a large variety of problems. The fact that both unrelated and independent programs give results to the $10\\%$ level of precision gives us confidence on both packages. The ZHS program is also suitable for simulations of experiments that use ice or water and the \\v Cerenkov imaging technique, as long as muons are not important. It can also be extended to other media. In addition, ZHS is considerably faster than the GEANT code and can reach much higher energies, making it adequate for high energy neutrino experiments.  Small discrepancies that remain between ZHS and GEANT are possibly due to minor differences which can be attributed to physical uncertainties in the processes involved. As a result all previously published calculations based on the ZHS Monte Carlo are expected to be correct to the same degree of accuracy, modulo other uncertainties which were discussed in the original papers."
    },
    "0206/astro-ph0206391.txt": {
        "abstract": "The combination of mid-infrared data from the MSX satellite mission and ground-based near-infrared photometry is used to characterise the properties of the mid-infrared population of the Galactic plane.  The colours of the youngest sources still heavily embedded within their natal molecular clouds are in general different from evolved stars shrouded within their own dust shells. Our main motivation is to use MSX for an unbiased search for a large ($\\sim$ 1000) sample of massive young stellar objects (MYSOs). A simple analysis shows that the MSX point source catalogue should contain most of the MYSOs within our Galaxy. We develop colour selection criteria using combined near- and mid-infrared data for MYSOs, which produces a list of 3071 objects, excluding the galactic centre region.  The programme of follow-up observations already underway to separate the MYSOs from compact H~II regions and other remaining objects is briefly described. We also show that these data can be used, just as IRAS data has been previously, to provide a separation between evolved stars with carbon rich and oxygen rich dust.  These data may also be used to search for evidence of dust around normal main sequence stars, such as low mass pre-main sequence stars or the Vega-excess class of objects where debris disks are presumed to remain from the planet formation process.  We discuss the accuracy and completeness of the MSX point source catalogue, and show that the errors present tend to be of a kind that is not significant for the main stellar populations we discuss in this paper. ",
        "introduction": "Objects embedded in dust play an important role throughout stellar evolution. Stars are born out of dusty molecular clouds whilst in the late stages of evolution they will invariably go through phases where they generate their own dust during heavy mass-loss. These objects emit most prodigiously in the infrared region of the spectrum with dust temperatures usually in the range of 30-300 K. To study these objects there is an undeniable need for unbiased samples and these usually have to originate from the infrared where most of their bolometric luminosity emerges.  Unbiased infrared surveys started with the rocket-borne AFGL mid-infrared survey of Price (1977) and then the first far-infrared satellite survey came with the IRAS mission. Many studies have been made using IRAS colours to select and classify objects of a variety of types (see Beichman 1987).  These can be split roughly into studies of evolved stars and studies of young stars.  In the former category, Van der Veen \\& Habing (1988) and Walker \\& Cohen (1988) used 12, 25 and 60$\\mu$m colours to classify the stars in the IRAS point source catalogue (IRAS PSC). These were mostly Asymptotic Giant Branch (AGB) stars that could be separated into those that were either oxygen-rich or carbon-rich, although first ascent Red Giant Branch (RGB) stars also have dust excesses (Plets et al.\\ 1997; Jura 1999). Searches for evolved stars where mass-loss had stopped and the dust shells have become detached were made by Oudmaijer et al.\\ (1992) for post-AGB stars and for planetary nebulae (e.g. Pottasch et al.\\ 1988).  Another example of the potential of IRAS for stellar astronomy was the discovery that some main sequence stars showed excess dust emission due to a remnant disk, now known as Vega-type or Vega-excess stars (Aumann 1985).  The other major area that benefited from IRAS was the study of the young stellar population in our galaxy.  IRAS colour-colour diagrams were used by Hughes \\& MacLeod (1988) to separate H~II regions from planetary nebulae and reflection nebulae for example. Prusti, Adorf \\& Meurs (1992) used more sophisticated techniques to extract low-mass young stellar object candidates from the IRAS PSC.  Our main motivation is to build a large unbiased sample of massive young stellar objects (MYSOs), also known after the prototype in Orion as BN objects (see Henning et al.\\ 1984). By this we mean a luminous ($\\gtappeq10^{4}$ \\Lsolar\\ or early B star), embedded source that has not yet reached the stage when the emergent Lyman continuum is sufficient to ionize the surrounding interstellar medium to form an H~II region. Due to the short contraction timescale compared to the collapse timescale it is thought that these objects are already core H-burning. However, it is also likely that accretion is still ongoing. They usually possess strong ionized stellar winds (e.g. Bunn, Hoare \\& Drew 1995) and drive powerful bipolar molecular outflows (Lada 1985). The 30 or so currently well-known MYSOs form a heterogeneous sample (Wynn-Williams 1982; Henning et al.\\ 1984), often found serendipitously, and may not be representative of the class as a whole.  There is an obvious need for a much larger well-selected sample.  It is reasonable at this point to ask how many MYSOs are in the entire galaxy. We can make a crude estimate by normalizing a stellar mass formation rate over a Salpeter IMF to yield the currently accepted star formation rate in the galaxy of about 6\\smpy\\ (G\\\"{u}sten \\& Mezger 1982). We then simply used the Kelvin-Helmholtz timescale for contraction to the main sequence using mass-radius and mass-luminosity relations for main sequence stars.  Although not strictly correct this does give appropriate timescales for the YSO phases that we are interested in, namely about 10$^{4}$ years for the most massive O star and 10$^{5}$ years for the intermediate mass Herbig Be type stars. This predicts about 1000 massive (15-100\\Msolar) YSOs in the galaxy, which is well over an order of magnitude larger than the current sample. A sample of this size is required when studying the properties of MYSOs as functions of mass, age, metallicity, etc. Hence, the task in reality is close to that of finding every MYSO in the galaxy.  Previous attempts at systematic searches for MYSOs have been made using IRAS data.  Campbell, Persson \\& Matthews (1989) in a series of papers followed-up 400 colour-selected IRAS sources that were also bright, unconfused and not identified with a known source. The problem with any infrared colour-selection procedure is that there are other types of source that have very similar infrared colours. Basically any heat source inside an optically thick dust cloud will produce an emergent spectral energy distribution that depends mostly on the optical depth of the cloud.  The population with most similar infrared colours to MYSOs are compact H~II regions, for which they are the progenitors. Both are are still deeply embedded in dense molecular clouds. Young compact, planetary nebulae and very dusty evolved stars can also have similar infrared colours. Campbell et al.\\ used single-element large aperture (8~arcsecond) near-infrared photometry to observe their IRAS candidates. They found a total of 115 candidate YSOs that were bright and red near-infrared sources, with excess emission in the $K$-band presumed to be from hot circumstellar dust. These candidates have luminosities in the range 10-10$^4$\\Lsolar\\ however so only just meet our criteria for a genuinely massive YSO.  Recent spectroscopy presented by Ishii et al.\\ (2001) suggests that some are not YSOs since the HI Br$\\gamma$ equivalent width is much closer to that seen in H~II regions (eg Lumsden \\& Puxley 1996) or PN (eg Lumsden, Puxley \\& Hoare 2001), than massive YSOs (eg Porter, Drew \\& Lumsden 1998).  This highlights the importance of obtaining sufficient follow up data to confirm the nature of any candidate MYSO.  Chan, Henning \\& Schreyer (1996) used similar selection criteria to Campbell et al.\\ and derived a list of 254 MYSO candidates. However, no attempt was made to eliminate compact H~II regions from the sample. Indeed over 100 of the objects are listed as being strong radio sources, much brighter than the weak stellar wind emission from nearby MYSOs (S$_{\\nu}<$10 mJy) (e.g. Henning, Pfau \\& Altenhoff 1990; Tofani et al.\\ 1995), and are therefore very likely to be H~II regions. Palla et al.\\ (1991) put together a colour-selected bright IRAS sample of 260 objects that avoided known H~II regions from single-dish radio surveys.  Of course, any infrared colour-selected sample of compact H~II regions by the converse will also contain large numbers of genuine MYSOs. Wood \\& Churchwell (1989a) developed 12, 25 and 60\\mic\\ colour-selection criteria based on ultra-compact H~II regions detected in their VLA survey. They applied this to the whole IRAS PSC and found 1646 embedded massive stars. Codella, Felli \\& Natale (1994) found that much more evolved, diffuse H~II regions as well as compact H~II regions satisfy the Wood \\& Churchwell criteria and Ramesh \\& Shridharan (1997) found evidence for significant contamination of the Wood \\& Churchwell sample by cores heated by less massive stars. Kurtz, Churchwell \\& Wood (1994) carried out a VLA survey of 59 bright IRAS sources satisfying the Wood \\& Churchwell colour criteria. They found that 80\\% had compact radio emission and are thus compact H~II regions, although some of these are weak enough to be MYSOs. Indeed several well-known MYSOs were recovered in their study. Sridharan et al.\\ (2002) selected a sample of 69 bright, northern MYSOs that satisfied the Wood \\& Churchwell criteria as well as having dense gas traced via CS emission and undetected in single-dish radio continuum surveys. The latter condition was to avoid complexes and H~II regions, although VLA follow-ups showed that nearly one third of their sample were still compact H~II regions. Their sample was deliberately biased towards finding isolated MYSOs for detailed high resolution studies.  Walsh et al.\\ (1998) carried out high resolution radio observations of IRAS sources originally selected in the same way as Wood \\& Churchwell, but mostly known to emit either radio continuum or methanol maser emission from previous single dish observations. Methanol masers have been known to trace massive star forming regions for some time (Menten 1991; Caswell et al.\\ 1995). High fractions of IRAS selected sources and well known MYSOs are maser sources, not only methanol, but H$_{2}$O and OH as well. However, unfortunately any one masing species cannot be guaranteed to be present in every MYSO. Indeed a recent search for methanol emission in a sample of well-known MYSOs found that none of them were detected (Gibb, Hoare \\& Minier, in preparation). So although there are great advantages in using strong, unobscured radio line emission techniques for galactic searches they will not be unbiased. Similarly, radio continuum emission from MYSOs is simply too weak to allow continuum mapping to be used in providing an unobscured search method for these objects.  Clearly therefore any unbiased search must start from infrared surveys. Preferably these should be at far-infrared wavelengths where the emission from the regions peaks and the extinction across the galaxy is negligible. Purely near-infrared surveys would be biased towards nearby and less heavily embedded (more evolved) sources due to the high extinction in the near-infrared. However, the concentrations of massive stars within $\\sim\\pm1^\\circ$ of the Galactic plane (Reed 2000; Becker et al.\\ 1994) means there is considerable source confusion in the IRAS survey due to its large beam (45~arcseconds $\\times$ 240~arcseconds at 12\\micron).  Recently the MSX satellite carried out a much higher spatial resolution survey (18~arcseconds) of the Galactic plane ($|b|<5^\\circ$) at 8, 12, 14 and 21$\\mu$m, completely superseding IRAS. This provides the starting point for our unbiased search for MYSOs.  The aim of this paper is to demonstrate that this mid-infrared data does provide a valuable means of locating these MYSOs, and to show how the various Galactic populations that emit in the mid-infrared can be characterised from such data. ",
        "conclusions": "We have shown that the combination of the MSX and 2MASS plus literature near-infrared data is a powerful diagnostic for detecting and classifying many kinds of embedded source with the Galactic Plane (in agreement with the earlier results of Egan 1999).  In particular, the combination provides a good separation between young embedded stars and evolved stars.  Evolved stars can easily be separated into those which have oxygen-rich dust and those which have carbon-rich dust, in a similar fashion to that first proposed by Epchtein et al.\\ (1987). The data are also likely to be a valuable tool in detecting normal stars with excess emission due to dust.  Although most of these are probably pre-main sequence stars, some may be systems such as Vega, where the dust emission is believed to arise in the debris disk created during the planet formation process.  We have developed colour selection criteria that will deliver a sample of about 2000 objects containing most of the massive young stellar objects and compact HII regions in our Galaxy. This number is in line with crude estimates of how many of these objects there should be in the Galaxy based on the IMF and current star formation rate. A programme of ground-based follow-up observations of these objects is already underway to confirm their identity and begin their detailed characterisation. This will lead to the first large ($\\approx$ 1000 objects) and well-selected sample of both MYSOs and compact H~II regions. Such an unbiased sample will be invaluable in the future study of massive star formation."
    },
    "0206/astro-ph0206333_arXiv.txt": {
        "abstract": "{ We have imaged five compact high--velocity clouds in \\hi with arcmin angular resolution and \\kms~spectral resolution using the Westerbork Synthesis Radio Telescope.  These CHVCs have a characteristic morphology, consisting of one or more quiescent, low--dispersion compact cores embedded in a diffuse warm halo.  The compact cores can be unambiguously identified with the cool neutral medium of condensed atomic hydrogen, since their linewidths are significantly narrower than the thermal linewidth of the warm neutral medium. Because of the limited sensitivity to diffuse emission inherent to interferometric data, the warm medium is not directly detected in the WSRT observations.  Supplementary total--power data, which is fully sensitive to both the cool and warm components of H\\,{\\sc i}, is available for comparison for all the sources, albeit with angular resolutions that vary from 3$^\\prime$ to 36$^\\prime$. The fractional \\hi flux in compact CNM components varies from 4\\% to 16\\% in our sample.  All objects have at least one local peak in the CNM column density which exceeds about $10^{19}\\rm\\;cm^{-2}$ when observed with arcmin resolution. It is plausible that a peak column density of 1--2$\\times10^{19}\\rm\\;cm^{-2}$ is a prerequisite for the long--term survival of these sources.  One object in our sample, CHVC\\,120$-$20$-$443 (Davies' cloud), lies in close projected proximity to the disk of M\\,31. This object is characterized by exceptionally broad linewidths in its CNM concentrations, more than 5 times greater than the median value found in the 13 CHVCs studied to date at comparable resolution. These CNM concentrations lie in an arc on the edge of the source facing the M\\,31 disk. The diffuse \\hi component of this source, seen in total--power data from the NRAO 140--foot telescope, has a positional offset in the direction of the M\\,31 disk. All of these attributes suggest that CHVC\\,120$-$20$-$443 is in a different evolutionary state than most of the other CHVCs which have been studied. Similarly broad CNM linewidths have only been detected in one other cloud, CHVC\\,110.6$-$07.0$-$466 (Wakker \\& Schwarz \\cite{wakker91b}) which also lies in the Local Group barycenter direction and has the most extreme radial velocity known.  A distinct possibility for Davies' cloud seems to be physical interaction of some type with M\\,31. The most likely form of this interaction might be the ram--pressure or tidal--stripping by either one of M\\,31's visible dwarf companions, M\\,32 or NGC\\,205, or else by a dark companion with an associated \\hi condensation. The compact objects located in the direction of the Local Group barycenter have an important role to play in constraining the Local Group hypothesis for the deployment of CHVCs. ",
        "introduction": "Although high--velocity clouds have been studied extensively since their discovery in~1963 by Muller \\ea~(\\cite{muller63}), there is still no consensus on the origin and physical properties of these objects. The clouds, for which most of the observations have been done in the \\hi 21--cm emission line, have velocities in excess of those allowed by Galactic rotation.  Most of the physical properties like size, mass, and gas density depend sensitively on the distances of the clouds: these distances are still unknown, except in a few cases.  The Magellanic Stream represents tidal debris originating in the gravitational interaction of the Large and Small Magellanic Clouds with our Galaxy (see Putman \\&~Gibson, \\cite{putman99}).  The Stream is therefore likely located at a distance of about 50 kpc.  Other high--velocity features with constrained distances are a few large complexes, extending over some tens of square degrees.  One of these, Complex~A, has been found from absorption--line observations (van Woerden et al.  \\cite{woerden99}, Wakker \\cite{wakker01}) to lie within the distance range of $8 < d < 10\\rm\\;kpc$.  Wakker \\&~van Woerden (\\cite{wakker97}) and Wakker et al.  (\\cite{wakker99}) have given recent reviews of the high--velocity cloud phenomenon.  During the past several years, there has been a renewed interest in the possibility that many high--velocity clouds are scattered throughout the Local Group. This hypothesis has been considered by many authors over the past decades: although some of the earlier references now appear somewat outdated, several early studies seem to have been particularly presentient: these include the work of Verschuur (\\cite{verschuur}), who discussed high--velocity clouds as protogalactic material scattered throughout the Local Group, and the work of Eichler (\\cite{eichler}) and Einasto et al.  (\\cite{einasto}), who viewed high--velocity clouds as carriers of dark matter, also scattered throughout the Local Group and available for merger with the larger systems.  Cosmological simulations intended to represent the evolution of the Local Group now predict a much higher number of dark--matter satellites around our Galaxy and Andromeda than the number of observed dwarf galaxies (Klypin et al.  \\cite{klypin99}, Moore et al. \\cite{moore99}).  Although there are several possible solutions to this problem, one is that the missing dark matter satellites should not only be identified with dwarf galaxies, but also with the high--velocity clouds.  These objects would have a very low star--formation rate, consistent with the non--detection of stars or of emission from dust or molecules associated with pre--stellar conditions.  Whereas Blitz et al. (\\cite{blitz99}) considered the properties of a general high--velocity cloud catalog in search of evidence for this hypothesis, Braun \\&~Burton (\\cite{braun99}) restricted their study to the compact and isolated ones, the so-called CHVCs.  These objects are isolated in the sense that they are not connected to extended emission features at a level of \\NH$ = 1.5\\times10^{18}$\\,cm$^{-2}$.  Such isolated objects turn out to be very compact, having a median angular size of less than $1^\\circ$.  The signature of these small and compact clouds in the Leiden/Dwingeloo survey (LDS, Hartmann \\&~Burton \\cite{hartmann97}) is indistinguishable from that of a nearby dwarf galaxy.  If the high--velocity clouds are the baryonic counterparts of low--mass dark--matter halos, then the subset of compact and isolated objects would be the most likely candidates for clouds at substantial distances, as yet undistorted by tidal-- and ram--pressure stripping.  The visual search for CHVCs of Braun \\&~Burton (\\cite{braun99}) in the LDS data has been extended by de\\,Heij et al. (\\cite{deheij02a}), with a fully automated algorithm.  The same algorithm was used to isolate the CHVC population in the southern hemisphere from the HIPASS data by Putman et al. (\\cite{putman02}).  The velocity dispersion of these compact and isolated clouds is the lowest in the Local Group Standard of rest system, lending some support to the idea that they are located throughout the Local Group.  More importantly, self--consistent modeling of \\hi bound to a dark--matter mini--halo population in the Local Group potential carried out by de\\,Heij et al. (\\cite{deheij02b}) gives support to this scenario. Critical aspects of this modeling are the realistic treatment of the effects of foreground obscuration by the \\hi of our Galaxy, and the account taken of the limited resolution and sensitivity of the existing survey data.  Due to its limited spatial resolution of~$36^\\prime$ FWHM, the Leiden/Dwingeloo survey is not an ideal basis for the study of the internal \\hi properties of the compact, high--velocity clouds.  Braun \\&~Burton (\\cite{braun00}) obtained high--resolution WSRT observations of six of these clouds. Other than the work by Braun \\& Burton and that reported here, only two CHVCs had previously been imaged at high resolution, by Wakker \\& Schwarz (\\cite{wakker91b}).  The synthesis observations reveal a characteristic morphology in which one or more compact cores are embedded in a diffuse halo, confirming the results from single--dish work done on large telescopes at moderately high angular resolution, notably in the earlier work done on the NRAO 300--foot telescope, whose FWHM beam subtended 10 arcminutes at $\\lambda21$ cm, by Giovanelli et al. (\\cite{giovanelli}).  The narrow line widths characteristic of most core components seen at arcminute resolution in the syntheis data allow unambiguous identification of these with the cool condensed phase of HI, the CNM, with kinematic temperatures near~$100\\rm\\;K$.  One of the CHVCs observed by Braun \\&~Burton (\\cite{braun00}), CHVC\\,125+41$-$207, showed several opaque clumps with some of the narrowest \\hi emission lines ever observed, with intrinsic FWHM of no more than~$2\\rm\\;km\\;s^{-1}$ and brightness temperature of $75\\rm\\;K$.  From a comparison of column and volume density for this object, Braun \\&~Burton estimate a distance in the range 0.5 to 1\\,Mpc.  In addition, several of the compact cores show systematic velocity gradients along the major axis of their elliptical extent. Some of these are well--fit by circular rotation in a flattened disk system. The apparent rotation velocities imply dark--matter masses of about 10$^8$~M$_\\odot$ and dark--to--visible mass ratios of 10\\,--\\,50 or more.  The cores of the multi--core objects show relative velocities as large as~$70\\rm\\;km\\;s^{-1}$ on 30~arcmin scales, also implying either an extremely short dynamically lifetime or a high dark--to--visible mass ratio.  In this paper, we extend the high--resolution study of CHVCs by imaging an additional five clouds with the WSRT.  Our discussion is organized as follows.  We begin by describing the method of sample selection in \\S\\,\\ref{sec:sample}, proceed with a description of the newly acquired observations in \\S\\,\\ref{sec:observations}, continue with a presentation of the images in \\S\\,\\ref{sec:presentation}, and conclude with discussion of our results in \\S\\,\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Cool and warm neutral media}  Wolfire~et al.\\ (\\cite{wolfire95}), following on the earlier treatments by Field~et al. (\\cite{field}) and Draine (\\cite{draine}), have shown that diffuse \\hi clouds in thermodynamic equilibrium might have a two--phase temperature structure.  The two components, a cool one (CNM) with temperatures around~$100\\rm\\;K$, and a warm one (WNM) with temperatures around~$10^4\\rm\\;K$, can coexist in pressure equilibrium for thermal pressures, $P/k_B$, in the range of about 100 -- 2000~cm$^{-3}$\\,K.  The calculations presented in Wolfire et al.\\ (\\cite{wolfire95}) have been supplemented with new ones appropriate for a low metallicity population of CHVCs residing at significant distances in the Local Group environment. The results of these new calculations are shown in Fig.~13 of Braun \\& Burton (\\cite{braun00}), where equilibrium solutions are given for clouds with a shielding column density of~1 and $10\\times10^{19}\\rm\\;cm^{-2}$, a metallicity of 0.1~solar, and a dust--to--gas mass ratio of 0.1 times the value in the solar neighborhood.  The velocity FWHM of \\hi clouds with kinetic temperatures of 100 K and~$10^4$ K are 2.4 \\kms~and 24 \\kms, respectively. As shown in Table~\\ref{table:gauss}, the values observed in the high--colum--density cores detected in the WSRT observations of our sample vary between FWHM of about 4 and 30~\\kms. No new example of ultra--narrrow \\hi lines was detected, such as the 2~\\kms~FWHM features seen in CHVC\\,125+41$-$207 by Braun \\& Burton (\\cite{braun00}).  The median linewidth of the material being discussed here amounts to about 6~\\kms~FWHM, a width comparable to that seen at high resolution in both HVC and CHVC studies (Wakker \\& Schwarz \\cite{wakker91}, Braun \\& Burton \\cite{braun00}) as well as in nearby external spiral galaxies (Braun \\cite{braun97}). This is somewhat broader than expected for the thermal linewidth of a 100~K gas, suggesting that either some form of ordered or random internal motions are present, or that the available resolution does not adequately account for line--of--sight blending of separate components.  The alternative, namely that the typical kinetic temperature is actually about 650~K, seems to be ruled out by observations of \\hi absorption in HVCs, as well as in the Galaxy and external galaxies, which reveal spin temperatures of 50 to 175~K in all cases (see Wakker et al.  \\cite{wakker91c}, Braun \\& Walterbos \\cite{braun92}, and Braun \\cite{braun95}).  The large velocity widths, of 25 to 30~\\kms~FWHM, found in the condensations along the eastern rim of CHVC 120$-$20$-$443 are difficult to interpret in this context. Of the thirteen CHVCs which have currently been subjected to arcmin resolution synthesis imaging, only one other instance of broad CNM linewidths has been observed, namely in CHVC\\,110.6$-$07.0$-$466 (Hulsbosch's cloud) as imaged by Wakker \\& Schwarz (\\cite{wakker91b}). Several other cases of broad widths are not relevant in this context, because they could be unambiguously attributed to line--of--sight overlap of components at different velocities.  As noted above in \\S\\,\\ref{sec:presentation}, it is not clear whether the linewidths in this feature are intrinsic or whether the velocity field becomes systematically double--valued at this location. We will return to this issue in a following subsection.  In general, it seems clear that it is predominantly the CNM which is detected in the WSRT images for objects of $0\\fdg5$ or more in size. The smoothly distributed WNM can not be readily detected in the synthesis data.  The fractional flux of CNM in the five objects studied here varies from about 4\\% to 16\\%. This fraction varied from less than 1\\% to more than 50\\% in the Braun \\& Burton (\\cite{braun00}) sample, which spanned a larger range of source properties. It is noteworthy that in all CHVCs studied to date with high spatial resolution there has been at least a marginal detection of the CNM. Every one of the thirteen CHVCs studied to date has at least one local peak in the CNM column density which exceeds about $10^{19}\\rm\\;cm^{-2}$ when observed with arcmin resolution. The accompanying diffuse WNM halo reaches comparable peak column densities, of about 1 or 2$\\times10^{19}\\rm\\;cm^{-2}$, external to these peaks (Burton et al. \\cite{burton01}).  It is conceivable that a WNM halo column density of 1 or 2$\\times10^{19}\\rm\\;cm^{-2}$ is a prerequisite for the long--term survival of these sources. It may be no coincidence that Maloney (\\cite{maloney93}) and Corbelli \\& Salpeter (\\cite{corbelli93}) estimated this value as the critical column density needed to prevent complete ionization of the \\hi when exposed to the estimated extragalactic ionizing radiation field relevant for free--floating objects in the Local Group.  \\subsection{Velocity offsets of cool cores and warm halos}  In the case of CHVC\\,186$+$19$-$114, it was possible to make a detailed comparison of line profiles as measured in CNM cores using the WSRT with the sum of the CNM and WNM emission detected in the 3--arcmin beam of the Arecibo telescope. The CNM spectra show narrower intrinsic widths as well as some local differences in the centroid velocity, while the Arecibo spectra display broad--linewidth tails (consistent with a 10$^4$~K thermal component) and much less dramatic variation in the profile shape and centroid. Given the dominant role of the WNM, accounting for about 84\\% of the total \\hi flux in this source, these differences are not surprising. Although the WSRT and Arecibo velocity centroids often agree, there are a few isolated locations where the CNM component is offset from the total \\hi centroid by a few \\kms. If a systematic velocity offset had been apparent between the WSRT and Arecibo spectra, it might have been an indication for an external perturbation of the source.  Br\\\"uns et al. (\\cite{bruens01}), who have observed the interesting object CHVC\\,125$+$41$-$207 with the 100--m Effelsberg telescope, argue that there is a systematic velocity offset between a narrow and broad component of the \\hi emission in that source. Their conclusion is based on Gaussian decompositions of the slightly asymmetric line profiles in the Effelsberg spectra. The decompositions result in two components; one of about 5 \\kms~and the other of 12 \\kms~FWHM. It is difficult to assess the physical relevance of these decomposition results, since at large distances from the CHVC centroid a single Gaussian of about 20~\\kms~is found to suffice in fitting the Effelsberg spectra well, while within the CHVC centroid, the WSRT data of Braun \\& Burton (\\cite{braun00}) for this object show non--Gaussian CNM line profiles of only 2 to 4~\\kms~FWHM.  Given the intrinsic non--uniqueness of Gaussian decomposition when applied to non--Gaussian line profiles, it seems questionable whether the 5 and 12 \\kms~FWHM Gaussian--fit components refer to physical systems at all.  If a systematic offset of the CNM and WNM velocities were present, then this might indicate that the halo kinematics is perturbed by an external force, which has not yet perturbed the central core of the cloud.  The gravitational tidal field of either the Galaxy or M\\,31 is a candidate for such a differential force. Another possibility is the ram--pressure exerted on the cloud as it moves through an external medium.  Given the substantial differences in sound--crossing times of the cores relative to the halo, a significant time delay in the response might result.  \\subsection{The particular interest of CHVC 120$-$20$-$443}  In his 1975 paper, Davies considered two possible interpretations of this cloud.  Given its proximity to M\\,31 on the sky, it might be located at a comparable distance, with a projected separation of only~$18\\rm\\;kpc$.  Since peak column densities are only a few times $10^{19}\\rm\\;cm^{-2}$, internal star formation is unlikely: our non--detection of stars in Palomar Sky Survey prints is no surprise. With only the visible baryonic mass, Davies concluded that the cloud is not gravitationally bound, and will double its size on a time scale of~$2.4\\times10^8$~years.  As an alternative possibility he considered that the cloud might be related in some way to the Magellanic Stream.  The closest approach of this feature to portions of the Stream is, however, about 30$^\\circ$ in angle and about 65~\\kms~in velocity, making such an association tenuous at best.  If the cloud were a part of the Magellanic Stream, its distance might be about $60\\rm\\;kpc$.  If there were no confining force except the self--gravity of the cloud, it would double its size in approximately $2\\times10^7\\rm\\;years$.  Our high--resolution imaging of CHVC\\,120$-$20$-$443 provides some insights into the possible origin of this object. As noted in \\S\\,\\ref{sec:presentation}, the high--column--density cores in this source are concentrated in a semi--circular rim along the eastern periphery, in the direction of the M\\,31 disk. Furthermore, exceptionally broad linewidths, of 25 to 30~\\kms~FWHM, are seen in this rim feature, while enhanced linewidths, amounting to 15 to 20~\\kms, are seen throughout the source. Of the thirteen CHVCs studied to date with arcmin resolution, only CHVC\\,110.6$-$07.0$-$466 has shown comparably broad linewidths in the CNM cores that are detected in interferometric data.  Median linewidths in the CNM cores of CHVCs imaged by Wakker \\& Schwarz \\cite{wakker91}, Braun \\& Burton \\cite{braun00}, and in this paper are only 6~\\kms. As noted previously, it is not yet clear whether the broad linewidths are intrinsic, or due to a large-scale geometric effect. One possibility might be a large physical extent along the line--of--sight.  Another curious circumstance is the large spatial offset of the brightest diffuse \\hi detected in the Green Bank 140--foot data toward the southeast of the CNM rim, as seen in Fig.~\\ref{fig:h120d}. All of these observations suggest that CHVC\\,120$-$20$-$443 is in a different evolutionary state than the other CHVCs which have been studied. Wakker \\& Schwarz (\\cite{wakker91b}) suggest a similarly different evolutionary state for CHVC\\,110.6$-$07.0$-$466.  A distinct possibility seems to be a physical interaction of some type with M\\,31.  It is interesting to speculate how an observer in M\\,31 would see CHVC\\,120$-$20$-$443 if it were at the relative distance of 18~kpc.  Given the properties of the cloud, the M\\,31 observer's perception of it could resemble the impression an earth-based observer has of the HVC Complexes~A or~C. For an observer located in the center of M\\,31, the cloud would extend over some~$30^\\circ$ on the sky.  Lower limits to the peak column densities that the observer would measure are determined by the ones measured in the WSRT observations, which have values of a few times $10^{19}\\rm\\;cm^{-2}$.  The WSRT observations show a filamentary structure with several embedded higher--density clumps.  The relative velocity of the object would be about 140~\\kms, given the M\\,31 systemic velocity of $-$300~\\kms. In order for this velocity to correspond to infall toward M\\,31 the object would have to be located beyond M\\,31, rather than between M\\,31 and the Galaxy.  From our vantage point in the Galaxy, the HVC Complex~A extends over about~$30^\\circ$ on the sky, while Complex~C extends over some~$70^\\circ$. Both have radial velocities of about $-$100~\\kms~in the Galactic Standard of Rest frame, and peak column densities of about $10^{19}\\rm\\;cm^{-2}$ as measured in the Leiden/Dwingeloo survey. Concerning distances, we note that Complex~A is well constrained to lie between 8 and 10~kpc (van Woerden et al. \\cite{woerden99}, Wakker \\cite{wakker01}), while only a few lower limits are available for Complex~C.  Although these clouds do not agree perfectly regarding their observable HI properties, they resemble each other sufficiently that it seems plausible to speculate about a similar physical origin.  Given the substantial projected distance of CHVC\\,120$-$20$-$443 from M\\,31, an origin in a galactic fountain within that galaxy seems unlikely.  In a galactic fountain, gas which is heated and ionised by supernova explosions rises to higher z--height, either buoyantly or driven by subsequent supernovae, where it finally condenses and returns in free fall back toward the galactic disk (see Shapiro \\& Field \\cite{shap76}, Bregman~\\cite{bregman80}). Simulations carried out by de\\,Avillez (\\cite{avillez00}) suggest that the height of this condensation process is at most several kpc above the stellar disk. CHVC\\,120$-$20$-$443 is located substantially further away from the stellar disk of M\\,31.  Because the driving force of a galactic fountain is provided by supernova explosions, which are concentrated in OB--associations, it is remarkable that only one such cloud would be seen.  The location of CHVC\\,120$-$20$-$443 is also not correlated with any region in M\\,31 of particularly active star formation (see Pellet et al. \\cite{pellet}), making this scenario appear unlikely.  A tidal origin for CHVC\\,120$-$20$-$443, related to either M\\,32 or NGC~205, is worth considering.  Ibata et al. (\\cite{ibata01}) have discovered a tidal stream of metal--rich stars extending several degrees toward the south of M\\,31.  They consider the dwarfs M\\,32 or NGC~205 as possibly responsible for the origin for the stream.  The angular extent of the stellar stream toward the south is comparable to the separation of CHVC\\,120$-$20$-$443 from the center of M\\,31 toward the north. Together these systems might trace portions of the same orbital path.  However, the measured radial velocity of the cloud is difficult to reconcile with those of the dwarfs. Both dwarfs have positive radial velocities with respect to~M\\,31 ($+155\\rm\\;km\\;s^{-1}$ in the case of M\\,32 and $+59\\rm\\;km\\;s^{-1}$ in the case of NGC~205) whereas the high--velocity cloud has a negative relative velocity of $-145\\rm\\;km\\;s^{-1}$.  According to the distances listed in Mateo (\\cite{mateo98}), NGC~205 is located beyond M\\,31.  Combined with its positive velocity with respect to M\\,31, it could be moving away from its peri--center passage.  During closest approach, the gas could have been stripped, either by ram--pressure stripping or by tidal disruption. However, the deceleration of the gas by some $200\\rm\\;km\\;s^{-1}$ would need to be accounted for. Realistic hydrodynamic simulations of such encounters might be illuminating.    Finally, the cloud could be part of a Local Group population of \\hi condensations within low--mass dark-matter halos, as described in the Local Group deployment model of CHVCs (Blitz et al. \\cite{blitz99}, Braun \\& Burton \\cite{braun99}). Analysis of the all--sky population of CHVCs performed by de\\,Heij et al. (\\cite{deheij02b}) has resulted in a self--consistent scenario whereby the observed CHVCs are part of a power--law distribution in baryonic mass (with slope $-1.7$) coupled to a steeper power--law (with slope $-$2) in dark mass. Only within the \\hi mass range of some 10$^{5.5}$ to 10$^7$ M$_\\odot$ are the objects stable against complete ionization by the intergalactic radiation field on the one hand (at low mass), and stable to internal star formation on the other (at high mass). The best--fitting simulated spatial distributions are centered on each of the Galaxy and M\\,31 with a spatial Gaussian dispersion of some 150~kpc. The majority of currently detected CHVCs belong to the relatively nearby swarm centered on the Galaxy. Only a small fraction of the M\\,31 sub--concentration of CHVCs is predicted to have been bright enough for detection in the current \\hi surveys. At the distance of M\\,31, CHVC\\,120$-$20$-$443 has an \\hi mass of about 10$^7$ M$_\\odot$, putting it at the high--mass end of the distribution.  If the projected separation with respect to M\\,31 is a measure for its real distance, then the cloud is sufficiently close to be strongly perturbed by the ram--pressure of its motion through a gaseous halo around M\\,31 (see de\\,Heij et al.  \\cite{deheij02b}). The observed extreme CNM linewidths in this object, and the significant displacement of the diffuse gas in the direction of M\\,31 with respect to the core components, may both be evidence for such an ongoing perturbation.  Of all of the CHVCs extracted by de\\,Heij et al. (\\cite{deheij02a}) from the LDS together with those found in the HIPASS material by Putman et al. (\\cite{putman02}) and comprising an all--sky sample, only six have a velocity more extreme than $|V_{\\rm LSR}|=400$ \\kms.  All of these objects have negative velocities, and all lie at northern declinations; they constitute the population of clouds often called VHVCs.  Arguments that this kinematic envelope is not an artifact of the observational parameters are given by de\\,Heij et al. (\\cite{deheij02b}). (The most extreme positive--velocity CHVC is the HIPASS object CHVC\\,258.2$-$23.9$+$359; the most extreme negative--velocity CHVC at southern declinations is CHVC\\,125.1$-$66.4$-$353.)  The most extreme--velocity CHVCs are the following, using the designation given by de\\,Heij et al. and, in parenthesis, the entry numbers from the catalogs of Wakker \\& van Woerden (\\cite{wakker91}), Braun \\& Burton (\\cite{braun99}), and de\\,Heij et al. (\\cite{deheij02a}): CHVC\\,103.4$-$40.1$-$414 (WW\\#491, deH\\#57), CHVC\\,107.7$-$29.7$-$429 (WW\\#437, BB\\#22, deH\\#59), CHVC\\,108.3$-$21.2$-$402 (WW\\#389, BB\\#23, deH\\#60), CHVC\\,110.6$-$07.0$-$466 (WW\\#318, BB\\#24, deH\\#61), CHVC\\,113.7$-$10.6$-$441 (WW\\#330, BB\\#25, deH\\#62), and Davies' cloud CHVC\\,120.2$-$20.0$-$444 (deH\\#68).  These CHVCs cluster near the direction of the barycenter of the Local Group, and are characteristically faint and small: they are likely to play an important role in the continuing discussion of the Local Group hypothesis.  The simulations of the Local Group hypothesis reported by de\\,Heij et al.  (\\cite{deheij02b}) support the prediction that a substantial number of additional CHVCs at extreme velocities will be found in the general direction of the Local Group barycenter, i.e. near M\\,31, when the sensitivity of the available \\hi survey data is improved.  The unusual properties of Davies' cloud may be revealed by other objects. Two of the extreme--velocity objects (both discovered by Hulsbosch, \\cite{hulsbosch}), namely CHVC\\,113.7$-$10.6$-$441 and CHVC\\,110.6$-$07.0$-$466, have been subject to synthesis imaging by Wakker \\& Schwarz (\\cite{wakker91b}).  It is interesting to note that Wakker \\& Schwarz state that the properties of these CHVCs differ considerably from the properties of the extended HVCs which they also partly imaged. CHVC\\,110.6$-$07.0$-$466 showed the same broad linewidth properties as we have found here for Davies' cloud.  It is plausible that the two objects have undergone a similar evolutionary experience.   \\subsection{Summary and conclusions}  We have imaged five CHVCs in \\hi with arcmin angular resolution and \\kms~spectral resolution using the Westerbork Synthesis Radio Telescope. These five images raise to 13 the number of CHVCs which have been subject to synthesis mapping, including the two compact objects studied by Wakker \\& Schwarz (\\cite{wakker91b}) and the six studied by Braun \\& Burton (\\cite{braun00}). These objects have a characteristic morphology, consisting of one or more quiescent, low--dispersion compact cores embedded in a diffuse warm halo.  The compact cores can be unambiguously identified with the cool neutral medium of condensed atomic hydrogen, since their linewidths are significantly narrower than the thermal linewidth of the warm neutral medium. Because of the limited sensitivity to diffuse emission inherent to interferometric data, the warm medium is not directly detected in the synthesis observations discussed here.  Supplementary total--power data, which is fully sensitive to both the cool and warm components of H\\,{\\sc i}, is available for all sources for comparison, although with angular resolutions that vary from 3$^\\prime$ to 36$^\\prime$. The fractional \\hi flux in compact CNM components varies from 4\\% to 16\\% in our sample.  All objects have at least one local peak in the CNM column density which exceeds about $10^{19}\\rm\\;cm^{-2}$ when observed with arcmin resolution. The accompanying diffuse WNM halo reaches comparable peak column densities, of about 1--2$\\times10^{19}\\rm\\;cm^{-2}$, external to these peaks (Burton et al. \\cite{burton01}).  It is conceivable that a WNM halo column density of 1--2$\\times10^{19}\\rm\\;cm^{-2}$ is a prerequisite for the long--term survival of these sources.  One object in our sample, CHVC\\,120$-$20$-$443 (Davies' cloud), lies in close projected proximity to the disk of M\\,31. This object is characterized by extremely broad linewidths in its CNM concentrations, which are 5 to 6 times broader than the median value found in the 13 objects studied to date at comparable resolutions. The CNM concentrations lie in an arc on the edge of the source facing the M\\,31 disk. The diffuse \\hi component of this source, seen in total--power data, has a large positional offset in the direction of the M\\,31 disk. All of these attributes suggest that CHVC\\,120$-$20$-$443 is in a very different evolutionary state than the other CHVCs which have been studied, with the possible exception of CHVC\\,110.6$-$07.0$-$466 (Hulsbosch's cloud), imaged by Wakker \\& Schwarz (\\cite{wakker91b}) and shown to also have broad CNM clumps. A distinct possibility seems to be a physical interaction of some type with M\\,31. The most likely form of this interaction might be ram--pressure or tidal--stripping by one of M\\,31's visible dwarf companions, M\\,32 or NGC~205, or by a dark companion with an associated \\hi condensation."
    },
    "0206/astro-ph0206105_arXiv.txt": {
        "abstract": "The Herbig Ae/Be star \\lkha\\ has been imaged at high angular resolution at a number of wavelengths in the near-infrared (from 1 $\\sim$ 3\\micronn) using the Keck~1 Telescope, and also observed in the mid-infrared (11.15\\micronn) using the U.C.~Berkeley Infrared Spatial Interferometer (ISI). The resolved circular disk with a central hole or cavity reported in \\citet{lk_nature} is confirmed. This is consistent with an almost face-on view (inclination $\\simle 35 ^\\circ$) onto a luminous pre- or early-main sequence object surrounded by a massive circumstellar disk. With a multiple-epoch study spanning almost four years, relative motion of the binary companion has been detected, together with evidence for changes in the brightness distribution of the central disk/star. The resolution of the \\lkha\\  disk by ISI mid-infrared interferometry constitutes the first such measurement of a young stellar object in this wavelength region. The angular size was found to increase only slowly from 1.6 to 11.15\\micronn, inconsistent with standard power-law temperature profiles usually encountered in the literature, supporting instead models with a hot inner cavity and relatively rapid transition to a cool or tenuous outer disk. The radius of the dust-free inner cavity is consistent with a model of sublimation of dust in equilibrium with the stellar radiation field. Measurements from interferometry have been combined with published photometry enabling an investigation of the energetics and fundamental properties of this prototypical system. ",
        "introduction": "Among the brightest young stellar objects in the infrared sky, \\lkha\\ is thought to be a transitional object which is on (or nearly on) the main sequence, but is still surrounded by a massive circumstellar disk. Following active accretion, hot stars are expected to pass through a brief phase in which the remnant accretion disk, which retains up to $\\sim$0.3 times the mass of the central star \\citep{Shu_90,Hollenbach_94}, is sculpted and eventually dissipated by the radiation and wind from the newborn star. The structure of these disks has been speculative, with uncertainty surrounding characteristic sizes and the existence of an inner cavity.  Images of young circumstellar disks in Orion\\citep{Odell_93, McG_O_96} and remnant or fossil disks at large distances from the star around ``Vega-type'' sources such as $\\beta$~Pic \\citep{Smith_84} have been seen in reflected light, or silhouetted against bright nebulosity. This work has attracted intense interest for, as first hypothesized by \\citet{kant}, it is likely that our own solar system grew out of such a flattened primordial nebula \\citep{Beckwith_96} labelled by Kant an ``Urnebel''.  The framework for understanding stellar formation is based on a rapidly-evolving theoretical and observational picture. Although the formation of circumstellar disks has long been favored, for the case of massive stars there have been historical problems with getting models which fit the shape of the spectral energy distribution (SED) \\citep{Hillenbrand_92} to account for the luminosities in the near-IR \\citep{Hartmann93}, and with forbidden line profiles not matching expectations \\citep{Bohm_Catala_94}. It was also found that the SED could be fitted by spherically symmetric shells \\citep{Miroshnichenko_97,Pezzuto_97} or composite shell-disk models \\citep{Miroshnichenko_99}.  Recent observational evidence has argued both for and against the presence of disks. Millimeter-wave interferometry has found evidence for large-scale rotation in Herbig~Ae stars \\citep{Mannings_Sargent97,Mannings_Sargent00}, albeit at much larger spatial scales than the expected sizes of the inner disks. However, similar studies of Herbig~Be stars \\citep{Fuente_01} conclude that the disks have dissipated at an early stage before the star becomes visible. Even restricting attention to just Herbig~Ae stars, HST coronographic observations of AB~Aur \\citep{Grady_99} show a more circularly symmetric structure consistent with disk viewed pole-on, contrasting with the mm observations.  Recent high-resolution imaging in the infrared has also been equivocal. Classical accretion disks have generally performed poorly in attempting to fit observations from the latest generation of long-baseline interferometers \\citep{IOTA_ABAur_99,Raphael_01,Akeson_00}. More directly, the expected source asymmetries from a population of disks, most of which must be inclined to the line of sight, were not indicated \\citep{Raphael_01}. However, of the three systems for which the hot inner regions have been well resolved by full interferometric imaging in the infrared, HK~Tauri~B \\citep{koresko98}, \\lkha\\ \\citep{lk_nature} and MWC~349 \\citep{mwc349}, clear disk morphologies (edge-on, face-on and edge-on respectively) have been established. In contrast to the near-IR where disks have proven larger than expected, \\citet{Hinz_01} failed to detect extended flux around a sample of Herbig Ae stars from mid-IR nulling interferometry.  Modellers are now re-thinking the simple thin/flared-disk geometries with power-law thermal profiles, now seen by many as inadequate in light of the measurements of Millan-Gabet and others, and promising new candidates for reconciling some of the discrepancies are emerging. One new model \\citep{Dullemond_01} involves an inner circumstellar cavity whose radius is set by thermal dust evaporation which adjoins a flaring disk geometry \\citep{Chiang_Goldreich_97} at larger radii. Such flared disks  may have thermal instabilities causing runaway local heating, puffed walls and self-shadowing of more distant regions \\citep{Dullemond_00}. The resultant geometry is capable of intercepting more of the stellar radiation, thus boosting its near-IR excess while maintaining excellent fits to the SED. This has been verified experimentally with new spectral data on four young stars \\citep{Natt_01}.  \\lkha was first identified as the source of illumination of the irregular reflection nebula NGC~1579 by Herbig (1956), who matched the H$\\alpha$ spectrum with a faint, deeply embedded star lying $\\sim$5 arcseconds within the border of a dark lane crossing the few-arcminute sized nebula. In the years since this discovery, it has become one of the most studied young stellar objects with observations spanning the spectrum. Despite this, pinning down many of its basic properties has proved elusive. The 800\\,pc distance established by \\citet{Herbig_71} was based on measurement of two nearby stars believed to be in association. Recently, \\citet{Stine_Oneal_98} made a strong case based on radio photometry that the system is much closer; suggesting instead that it is located in an extension of the Taurus-Auriga star formation complex at 160\\,pc.  Part of the difficulty has been the complexity of the highly anisotropic circumstellar environment, with at least four molecular clouds \\citep{Redman_86,Barsony_90} surrounding \\lkha\\ and its associated H~{\\small II} region, S222 \\citep{Herbig_56}. A deeply-embedded population of up to hundreds of probable protostars, pre-main-sequence low-mass stars, and brown dwarf candidates has been detected in the radio \\citep{Becker_White_88,Stine_Oneal_98} and infrared \\citep{Barsony_91,Aspin_Barsony_94}. Estimates of the visible extinction vary wildly in the literature by almost 10 magnitudes between values $A_v = $ 9.4 \\citep{Barsony_90} to 18.5\\,mag \\citep{Hou_97}.  This paper presents results from extremely high angular resolution imaging obtained with interferometric techniques. This has allowed observation of the hot, young, self-luminous disk surrounding \\lkha\\ in the near- and mid-infrared. The observational methods and the two different instruments used to obtain the data are described in Section~\\ref{obs}, with the results given in Section~\\ref{results}. The astrophysical interpretations are discussed in Section~\\ref{discussion}, while section Section~\\ref{conclusions} contains a summary of the important findings. ",
        "conclusions": "\\label{conclusions}  Near-infrared images of the Herbig Ae/Be star \\lkha\\ have been obtained from a multi-epoch study utilizing interferometry on the Keck~1 telescope to obtain information at the diffraction limit (tens of milli-arcseconds). The mid-infrared U.C. Berkeley Infrared Spatial Interferometer was also able to resolve the system at 11.15\\,$\\micron$. \\lkha\\ presents a resolved circular disk with a central hole or cavity. A relatively blue-spectrum binary companion 180\\,mas to the E-NE is also reported.  The morphology of the resolved limb-brightened ring is interpreted as a close to face-on viewing angle onto a geometrically thick torus or circumstellar disk with a hot inner wall facing the central star. The size of the cavity is consistent with the radiative equilibrium temperature for dust sublimation. A relatively slow increase in apparent diameter of the disk with observing wavelength over a decade from 1.2 to 11.15\\,$\\micron$ implies circumstellar material with a relatively compact density and/or thermal profile, arguing against classical power-law temperature profiles usually encountered in the literature.  Relative motion of the binary companion, together with simple assumptions on the masses and geometrical projection, favor a likely intermediate distance of around $\\sim$340\\,pc, although the earlier large (800\\,pc) and particularly small (160\\,pc) suggested distance scales cannot be ruled out.  Combining the interferometry with published photometry enables a decomposition of the spectral energy distribution into 3 component parts in the infrared. In addition to the 180\\,mas companion and the resolved disk, we report the probable isolation of a bright unresolved source at the center of the cavity in the disk which has been identified as the primary. This is likely an early B~star and the source of the photoionizing radiation driving the HII region. With more detailed radiative transfer modelling and studies at yet higher resolution from the latest generation of optical interferometers, this fascinating system is among the most promising candidates for in-depth study of the workings of a massive young star."
    },
    "0206/astro-ph0206111_arXiv.txt": {
        "abstract": "We report on the study of the mass-radius (M-R) relation and the radial oscillations of proto strange stars. For the quark matter we have employed the well known density dependent quark mass model and its very recent modification, the temperature and density dependent quark mass model. We find that the maximum mass the star can support increases significantly with the temperature of the star in this model which implies that transition to a black hole at the early stage of formation of the star is inhibited. As for the neutrinos, we find, contrary to the expectation that the M-R and oscillation frequencies are almost independent of the neutrino chemical potentials.\\\\ Subject headings: Strange stars - Oscillations ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206327_arXiv.txt": {
        "abstract": "{We calculate the two-, three- and (for the first time) four-point correlation functions of the \\emph{COBE}-DMR 4-year sky maps, and search for evidence of non-Gaussianity by comparing the data to Monte Carlo-simulations of the functions. The analysis is performed for the 53 and 90 GHz channels, and five linear combinations thereof. For each map, we simulate an ensemble of 10\\,000 Gaussian realizations based on an \\emph{a priori} best-fit scale-invariant cosmological power spectrum, the DMR beam pattern and instrument-specific noise properties. Each observed \\emph{COBE}-DMR map is compared to the ensemble using a simple $\\chi^2$ statistic, itself calibrated by simulations. In addition, under the assumption of Gaussian fluctuations, we find explicit expressions for the expected values of the four-point functions in terms of combinations of products of the two-point functions, then compare the observed four-point statistics to those predicted by the observed two-point function, using a redefined $\\chi^2$ statistic. Both tests accept the hypothesis that the DMR maps are consistent with Gaussian initial perturbations. ",
        "introduction": "The study of CMB temperature anisotropies and their statistical properties has become an important theme in modern cosmology. In its most conventional interpretation, the distribution of anisotropies reflects the properties of the universe approximately 300\\,000 years after the Big Bang, at the surface of last scattering. Thus, by measuring statistical quantities such as the angular power spectrum or the angular two-point correlation function, we can infer the values of many interesting cosmological parameters.  For both theoretical and practical purposes, it is convenient to expand the temperature anisotropy field into a sum of (complex) spherical harmonics: \\begin{equation} \\Delta T(\\theta, \\phi) = \\sum_{lm} a_{lm} \\: Y_{lm}(\\theta, \\phi) \\end{equation} The temperature perturbation field is said to be Gaussian distributed if each $a_{lm}$ follows an independent Gaussian probability distribution. The question of whether the observed temperature field is Gaussian or otherwise is of crucial importance for modern cosmology. From a scientific standpoint, most conventional inflationary models of structure formation predict a Gaussian temperature field, whereas scenarios which invoke topological defects to seed the large-scale structure predict a non-Gaussian distribution. Thus the statistical properties of the $a_{lm}$'s can be used to distinguish such models. Secondly, from an analysis point-of-view, most parameter estimation techniques -- usually based on the observed angular power spectrum -- assume Gaussianity, and may therefore be biased if the observed field is indeed non-Gaussian.  However, testing for non-Gaussianity is anything but trivial, and several qualitatively different tests are required in order to perform a complete analysis. At present the arsenal of available tests which have been applied to the \\emph{COBE}-DMR data consists of at least the following: bi- and trispectrum based analysis (Ferreira et al. \\cite{ferreira}; Magueijo \\cite{magueijo}; Sandvik \\& Magueijo \\cite{sandvik}; Komatsu et al. \\cite{komatsu}; Kunz et al. \\cite{kunz}), 3-point correlation function based tests (Kogut et al. \\cite{kogut}), methods utilizing wavelets (Cay\\'on et al. \\cite{cayon}; Barreiro et al. \\cite{barreiro}) and Minkowski functionals (Schmalzing \\& G\\'orski \\cite{schmalzing98}; Novikov et al. \\cite{novikov}). Indeed, there has been a small resurgence in interest in the possibility of non-Gaussian signals in the \\emph{COBE}-DMR maps as a consequence of the bispectrum work of Ferreira et al. (\\cite{ferreira}) and Magueijo (\\cite{magueijo}). These papers find non-Gaussian contributions using harmonic analyses at the 98\\% confidence limit, and although Banday et al. (\\cite{bandaya}) explain these tentative detections by appealing to the presence of a specific residual systematic artifact in the data, additional investigation is warranted.  In this paper we adopt $N$-point correlation functions as probes for non-Gaussianity. For a Gaussian field all odd $N$-point functions (such as the three-point function) have vanishing expectation values, while all even $N$-point functions can be reduced to expressions involving the two-point function. Thus, if the observed three-point function is significantly non-zero when compared to a Gaussian ensemble, its native distribution is probably non-Gaussian. Further, if the four-point function does not reduce into two-point functions, the same conclusion can be made.  The first part of this paper builds on ideas demonstrated in Kogut et al. (\\cite{kogut}) and Hinshaw et al. (\\cite{hinshaw}). We study the 4-year \\emph{COBE}-DMR sky maps, computing the two- and three-point functions, as has been performed previously, then proceeding to extend the analysis for the first time to the determination of several four-point functions. The definitions of these new functions are given in Sect.\\ \\ref{sec:definitions}.  In Sect.\\ \\ref{sec:measurements} we compute the various correlation functions for the four DMR channels and five linear combinations thereof. Next, we compute the same functions for 10\\,000 Monte Carlo simulated Gaussian maps, which are used as the basis of the various statistical tests of non-Gaussianity. Initially, we apply the $\\chi^2$ test as defined by Kogut et al. \\cite{kogut}, by comparing the observed data value to the distribution generated from the application of the $\\chi^2$ statistic for each map in the simulated ensemble.  Subsequently, in Sect.\\ \\ref{sec:reduction}, we provide expressions for the expected value of several four-point functions in terms of the two-point function, then explicitly compare the observed four-point functions to those predicted by the observed two-point function. We define a suitable $\\chi^2$ statistic in order to quantitatively measure the degree of deviation, once again calibrated by Monte Carlo-simulations. ",
        "conclusions": "By performing Monte Carlo-simulations we have studied the statistical properties of the \\emph{COBE}-DMR 53 GHz and 90 GHz channels. The basic ingredients for this analysis were various $N$-point correlation functions, and, in particular, four different four-point functions which have been presented for the first time. We have additionally taken advantage of a result from statistical theory, relating all even $N$-point functions to reductions in terms of the two-point function. This allowed us to define a test for Gaussianity in which the assumed power spectrum only plays a secondary role. This test could therefore prove better suited for situations in which we do not have access to the optimal power spectrum.  Comparison of the DMR $N$-point correlation functions with the Monte-Carlo ensemble indicates no evidence for possible non-Gaussian behavior, in agreement with the earlier analysis of Kogut et al. (\\cite{kogut}). Furthermore, the agreement between the observed DMR functions and the simulated ensembles also supports the validity of our model assumptions, namely that of a scale-invariant power law model for the anisotropies, and uncorrelated noise.  On the other hand, the excellent agreement between the simulated and the observed correlation functions poses an intriguing problem: tests of Gaussianity based on a harmonic analysis of the DMR data -- the bispectrum work of Ferreira et al. (\\cite{ferreira}) and trispectrum results of Kunz et al. (\\cite{kunz}) -- show compelling evidence for non-Gaussian features (although these have subsequently been associated with systematic artifacts in the DMR data by Banday et al. \\cite{bandaya}), while tests based on real-space high-order statistics such as those presented here do not. The resolution of such apparently contradictory results is most likely rather mundane: the source of the non-Gaussian signal was found to be strongly located at the multipole order $l=16$. Since the correlation functions are by definition (weighted) averages over the full multipole range, the reduced sensitivity to this type of non-Gaussian structure is certainly not unexpected."
    },
    "0206/astro-ph0206057_arXiv.txt": {
        "abstract": "{ We discuss the properties of the host galaxy and ring light distributions in the optical and near infrared bands for a sample of Polar Ring Galaxies (PRGs), presented in Paper I (Iodice et al. \\cite{paperI}). The goal of this work is to test different formation scenarios for PRGs, proposed by different authors in the last decades, by comparing their predictions with these new data. The strategy is twofold: {\\it i)} the integrated colors of the main components in these systems are compared with those of standard morphological galaxy types, to investigate whether differences in colors are caused by dust absorption or difference in stellar populations. We then derived an estimate of the stellar population ages in PRGs, which can be used to set constrains on the dynamical modeling and the time  evolution of these systems; {\\it ii)} we analyse the structural parameters of the host galaxy in order to understand whether this component is a standard early-type system as its morphology suggests, and the light distribution in the polar ring to measure its radial extension.\\\\ These observational results indicate that the global properties of PRGs are better explained by dissipative merging of disks with un-equal masses as proposed by Bekki (1998), rather than the accretion-or stripping-of gas by a pre-existing  early-type galaxy. ",
        "introduction": "Gravitational interactions and mergers affect the morphologies and dynamics of galaxies. Both the Hubble Space Telescope (HST) and the ESO Very Large Telescope (VLT) make it possible to observe the early universe and show that these processes dominate at higher redshift (Driver et al. \\cite{driver95}). Polar Ring Galaxies (PRGs) are considered one of the best example of remnants from galaxy interaction. These peculiar objects are composed by a central spheroidal component, the {\\it host galaxy}, surrounded by an outer {\\it ring}, made up by gas, stars and dust, in a nearly perpendicular plane to the equatorial one of the central galaxy (Whitmore et al. \\cite{PRC}). The origin of PRGs is still an open debate: they may be the result of an accretion phenomena (Toomre \\cite{Toomre77}; Shane \\cite{Shane80}; Schweizer et al. \\cite{Schweizer83}; Sackett \\cite{Sackett91}; Sparke \\cite{Sparke91}; Hibbard \\& Mihos \\cite{Hibbard95}; Reshetnikov and Sotnikova \\cite{Resh_Sot97}) or of major dissipative merger (Bekki \\cite{Bekki97}; \\cite{Bekki98}). Recent observations of several interacting galaxy pairs, of similar luminosities, display ring-like structures (e.g. NGC~7464/65, Li \\& Seaquist \\cite{Li94}; NGC~3808A,B and NGC~6285/86, Reshetnikov et al. \\cite{Resh96}).  Reshetnikov \\& Sotnikova (\\cite{Resh_Sot97}) studied the accretion scenario for the formation of PRGs using a smoothed-particle hydrodynamic simulations (SPH) in high speed encounters. They analyzed the different ring morphologies which were generated by the encounter of a gas-rich spiral with either an elliptical or an S0 galaxy. They followed the full history of the gas stripping: from the spiral galaxy outskirts to its capture by the early-type galaxy, on a parabolic encounter. The total amount of accreted gas by the early-type object is about $10\\%$ of the gas in the spiral galaxy, i.e. up to $10^{9}$ $M_{\\odot}$. The size of the polar ring is found to be related to the central mass (luminous + dark) concentration of the host galaxy.  If the mass is highly concentrated, i.e. the elliptical galaxy case, the ring forms at smaller radii; if the host galaxy has an extended massive halo, i.e. the S0 case, the ring average radius ($\\bar{R}$) can be as large as 30~kpc. This scenario can account for the formation of (quasi-)stable polar rings, whose radial extent is of the order of 10\\% of the ring extention.  A quite different approach to the formation of polar ring galaxies was recently proposed by Bekki (\\cite{Bekki98}). In this scenario, the polar ring results from a ``polar'' merger of two disk galaxies with unequal mass. The ``intruder'', on a polar orbit with respect to the ``victim'' disk, passes through it near its center: it is slowed down, and pulled back toward the victim, by strong dissipation which is caused by the interaction with the victim gaseous disk. Dissipation removes random kinetic energy from the gaseous component of the victim's disk, so that some gas can settle again into a disky configuration. The morphology of the merger remnants depends on the merging initial orbital parameters and the initial mass ratio of the two galaxies.  Bekki's scenario successfully reproduces the range of observed morphologies for polar ring galaxies, such as the existence of both wide and narrow rings, helical rings and double rings  (Whitmore \\cite{whitmore91}). Furthermore, this scenario would also explain the presence of wide and massive polar disk, as observed in NGC~4650A (Arnaboldi et al. \\cite{magda97}, Iodice et al. \\cite{4650aI}; Gallagher et al. \\cite{4650aG}). The two scenarios, accretion vs. dissipative mergers of disks, both predict the general features of PRGs: a structure-less appearance of the host galaxy, and the younger dustier appearance of the polar structure. But they differ on their predictions about structural parameters, age, baryonic mass and polar structure extension. Therefore, the two scenarios should be tested against the observed properties of both  wide and narrow PRGs, in particular their observed structural parameters, colors and total light. Given that PRGs contain a lot of dust, we must study their light distribution in the near-infrared (NIR), and determine the distribution of their evolved stellar population. To this aim, new NIR data for a sample of PRGs were collected and analysed in Iodice et al. \\cite{paperI}, hereafter Paper I. In this work we compare the integrated colors derived for the host galaxy and ring of each PRG in our sample (see Sec.5 in Paper I) with those of standard morphological galaxy types, in Sec.\\ref{nircol}, and compute an estimate of the stellar population ages, in Sec.\\ref{PR_age}. In Sec.~\\ref{model} and Sec.~\\ref{param_pr}, we perform a detailed analysis of the light distribution properties in the host galaxy and ring.  The new observational evidences obtained for this sample of PRGs are summarized in Sec.~\\ref{sum}, and conclusions are derived in Sec.~\\ref{conclu}. ",
        "conclusions": "In this work we have presented the analysis of the detailed photometric study for a sample of PRGs, selected from the Polar Ring Catalogue (PRC, Whitmore et al. \\cite{PRC}), based on new NIR observations presented in Iodice et al. \\cite{paperI} (Paper~I). We now wish to compare the properties predicted for PRGs in different formation scenarios, presented in  Sec.\\ref{intro} (see also Iodice et al. \\cite{4650aI}), against the global properties observed for the polar ring systems studied in this work.  For all PRGs of our sample, except for AM~2020-504, the published simulations of the accretion/stripping scenario are so far not able to predict \\begin{enumerate} \\item the main characteristics in the light and color distribution of the host galaxy, which make this component different from a standard S0 system; \\item the large values for the $\\Delta R / \\bar R$ ratio, which is related to the ring extension; \\item the large baryonic mass (stellar + gas), shown in Table~\\ref{tab_MB}, in the polar ring. \\end{enumerate}  The observed properties for AM~2020-504 (see Sec.7 in Paper~I) suggest that this polar ring may be the single case in our sample which may be formed through an accretion or gas-stripping involving an elliptical galaxy.  All the observed properties of the host galaxy and polar structure can be more easily explained by the dissipative merger scenario proposed by Bekki (\\cite{Bekki98}). In this scenario both the central S0-like system and ring component in a polar ring galaxy are simultaneously formed through a dissipative merger between two disk galaxies. The required constraints on the specific orbital configurations and gaseous dissipation in galaxy merging naturally explain the prevalence of S0-like systems among polar ring galaxies (e.g., Whitmore \\cite{whitmore91}) and the appreciably  larger amount of interstellar gas in PRGs (van Gorkom et al. \\cite{vgorkom87}, Arnaboldi et al. \\cite{magda97}, van Driel et al. \\cite{vdriel2000}). This scenario does predict peculiar characteristics for the host galaxy: the progenitor galaxy (the intruder) experiences both a heating of the disk (it puffs up) and energy dissipation. The energy dissipation leads to an higher increase of the mass density in the center, with respect to the unperturbed disk, which may develop a central small and nearly exponential bulge: this is very similar to what we have detected in nearly all PRGs of our sample. The evolutionary timescales of the merging process, which is about $10^9$ yr, is also consistent with the young age, predicted for PRGs in this work, both for the host galaxy ($1$ to $3$ Gyr) and polar structure ($\\sim 1$ Gyr). Furthermore,  the different morphologies observed for polar rings, such as narrow rings (e.g. ESO~415-G26, or in ARP~230) and wide disk-like structures with no central hole (e.g. NGC~4650A, see Iodice et al. \\cite{4650aI} and Gallagher et al. \\cite{4650aG}), are related to the orbital parameters of galaxy merging and the initial mass ratio of the two interacting galaxies.  An important constraint to the Bekki scenario is the small value of the relative velocities ($V \\sim 33$ km s$^{-1}$) that the two merging galaxies need to have to form PRGs: such velocities are more likely to occur in high redshift universe rather than nearby,  where bound group of galaxies are virialized and therefore their relative velocities are larger. Reshetnikov (\\cite{Resh97}) have found an increasing rate of detection for PRGs toward higher redshift: among all galaxy types, in the Hubble Deep Field (Williams et al.  \\cite{Williams95}) candidate polar ring galaxies are $\\sim 0.7\\%$, while in the local universe this is $\\sim 0.05\\%$ (Whitmore et al. \\cite{PRC}). Although uncertainties in the numerical treatment of gas dynamics and star formation still remain in the Bekki's approach, dissipative galaxy merging, with specific initial conditions, seems now a promising scenario to to explain the formation of Polar Ring Galaxies and their observational properties.   \\begin{table} \\caption[]{Mass of the stellar component in the host galaxy (second column) and the total baryonic mass in the polar structure (third column), which includes the mass of the stellar component and the mass of the gas in the form of neutral  (HI) and molecular hydrogen (HII). See Sec.7, Paper~I for additional references.} \\label{tab_MB} \\begin{tabular}{lcc} \\hline\\hline Object & $M_{star}$ (HG) & $M_{gas} + M_{star}$ (PR)\\\\ &  $10^9 M_{\\odot}$ & $10^9 M_{\\odot}$\\\\ \\hline ESO 415-G26 & 9 & 10\\\\ ARP 230 & 2 & 5\\\\ AM 2020-504& 6 & 5\\\\ ESO 603-G21& 2 & 10\\\\ \\hline \\end{tabular} \\end{table}"
    },
    "0206/astro-ph0206261_arXiv.txt": {
        "abstract": "The nonlinear coefficients defining the mean electromotive force ({\\em i.e.,} the nonlinear turbulent magnetic diffusion, the nonlinear effective velocity, the nonlinear $ \\kappa $-tensor, {\\em etc.}) are calculated for an anisotropic turbulence. A particular case of an anisotropic background turbulence ({\\em i.e.,} the turbulence with zero mean magnetic field) with one preferential direction is considered. It is shown that the toroidal and poloidal magnetic fields have different nonlinear turbulent magnetic diffusion coefficients. It is demonstrated that even for a homogeneous turbulence there is a nonlinear effective velocity which exhibits diamagnetic or paramagnetic properties depending on anisotropy of turbulence and level of magnetic fluctuations in the background turbulence. The diamagnetic velocity results in the field is pushed out from the regions with stronger mean magnetic field, while the paramagnetic velocity causes the magnetic field tends to be concentrated in the regions with stronger field. Analysis shows that an anisotropy of turbulence strongly affects the nonlinear turbulent magnetic diffusion, the nonlinear effective velocity and the nonlinear $ \\alpha $-effect. Two types of nonlinearities (algebraic and dynamic) are also discussed. The algebraic nonlinearity implies a nonlinear dependence of the mean electromotive force on the mean magnetic field. The dynamic nonlinearity is determined by a differential equation for the magnetic part of the $ \\alpha $-effect. It is shown that for the $ \\alpha \\Omega $ axisymmetric dynamo the algebraic nonlinearity alone (which includes the nonlinear $ \\alpha $-effect, the nonlinear turbulent magnetic diffusion, the nonlinear effective velocity, {\\em etc.}) cannot saturate the dynamo generated mean magnetic field while the combined effect of the algebraic and dynamic nonlinearities limits the mean magnetic field growth. ",
        "introduction": "Generation of magnetic fields by turbulent flow of conducting fluid is a fundamental problem which has a large number of applications in solar physics and astrophysics, geophysics and planetary physics, {\\em etc.} In recent time the problem of nonlinear mean-field magnetic dynamo is a subject of active discussions (see, {\\em e.g.,} \\cite{KA92,GD94,BB96,K99,F99,MS99,RK2000,KMRS2000,UB2000,VC01}). It was suggested in \\cite{VC92} that the quenching of the nonlinear $ \\alpha $-effect is very strong and causes a very weak saturated mean magnetic field. However, the later suggestion is in disagreement with observations of galactic and solar magnetic fields (see, {\\em e.g.,} \\cite{M78,P79,KR80,ZRS83,RSS88,S89}) and with numerical simulations (see, {\\em e.g.,} \\cite{GR95,BN95,BR96}).  Saturation of the dynamo generated mean magnetic field is caused by the nonlinear effects, {\\em i.e.,} by the back reaction of the mean magnetic field on the $ \\alpha $-effect, turbulent magnetic diffusion, differential rotation, {\\em etc.} The evolution of the mean magnetic field $ {\\bf B} $ is determined by equation \\begin{eqnarray} \\partial {\\bf B} / \\partial t = \\bec{\\nabla} \\times ({\\bf V} \\times {\\bf B} + \\bec{\\cal E} - \\eta \\bec{\\nabla} \\times {\\bf B}) \\;, \\label{A12} \\end{eqnarray} where $ {\\bf V} $ is a mean velocity ({\\em e.g.,} the differential rotation), $ \\eta $ is the magnetic diffusion due to the electrical conductivity of fluid. The mean electromotive force $ \\bec{\\cal E} = \\langle {\\bf u} \\times {\\bf b} \\rangle $ in an anisotropic turbulence is given by \\begin{eqnarray} {\\cal E}_{i} &=& \\alpha_{ij} B_{j} + ({\\bf V}^{\\rm eff} {\\bf \\times} {\\bf B})_{i} - \\eta_{ij} (\\bec{\\nabla} {\\bf \\times} {\\bf B})_{j} \\nonumber\\\\ && - \\kappa_{ijk} (\\partial \\hat B)_{jk} - [\\bec{\\delta} {\\bf \\times} (\\bec{\\nabla} {\\bf \\times} {\\bf B})]_{i} \\; \\label{A14} \\end{eqnarray} (see \\cite{R80,RKR00}), where $ (\\partial \\hat B)_{ij} = (1/2) (\\nabla_{i} B_{j} + \\nabla_{j} B_{i}) ,$ $ {\\bf u} $ and $ {\\bf b} $ are fluctuations of the velocity and magnetic field, respectively, angular brackets denote averaging over an ensemble of turbulent fluctuations, the tensors $ \\alpha_{ij} $ and $ \\eta_{ij} $ describe the $ \\alpha $-effect and turbulent magnetic diffusion, respectively, $ {\\bf V}^{\\rm eff} $ is the effective diamagnetic (or paramagnetic) velocity, $ \\kappa_{ijk} $ and $ \\bec{\\delta} $ describe a nontrivial behavior of the mean magnetic field in an anisotropic turbulence. Nonlinearities in the mean-field dynamo imply dependencies of the coefficients $ (\\alpha_{ij} , \\eta_{ij}, {\\bf V}^{\\rm eff}, $ {\\em etc.}) defining the mean electromotive force on the mean magnetic field. The $ \\alpha $-effect and the differential rotation are the sources of the generation of the mean magnetic field, while the turbulent magnetic diffusion and the $ \\kappa $-effect (which is determined by the tensor $ \\kappa_{ijk}) $ contribute to the dissipation of the mean magnetic field.  In spite the nonlinear $ \\alpha $-effect was under active study (see, {\\em e.g.,} \\cite{F99,RK2000}), the nonlinear turbulent magnetic diffusion, the nonlinear $ \\kappa $-effect, the nonlinear diamagnetic and paramagnetic effects, {\\em etc.} are poorly understood.  In the present paper we derived equations for the nonlinear turbulent magnetic diffusion, the nonlinear effective velocity, the nonlinear $ \\kappa $-effect, {\\em etc.} for an anisotropic turbulence. The obtained results for the nonlinear mean electromotive force are specified for an anisotropic background turbulence with one preferential direction. The background turbulence is the turbulence with zero mean magnetic field. We demonstrated that toroidal and poloidal magnetic fields have different nonlinear turbulent magnetic diffusion coefficients. It is shown that even for a homogeneous turbulence there is a nonlinear effective velocity which can be a diamagnetic or paramagnetic velocity depending on anisotropy of turbulence and level of magnetic fluctuations in the background turbulence. ",
        "conclusions": "In this study we calculated the nonlinear tensor of turbulent magnetic diffusion, the nonlinear $ \\kappa $-tensor, the nonlinear effective velocity, and other coefficients defining the mean electromotive force for an anisotropic turbulence. The obtained results were specified for an anisotropic background turbulence with one preferential direction. We found that the turbulent magnetic diffusion coefficients for the toroidal and poloidal magnetic fields are different. We demonstrated that even for a homogeneous turbulence there is the nonlinear effective velocity which can be a diamagnetic or paramagnetic velocity depending on anisotropy of turbulence and level of magnetic fluctuations in the background turbulence. The diamagnetic velocity implies that the field is pushed out from the regions with stronger mean magnetic field, while the paramagnetic velocity causes the magnetic field tends to be concentrated in the regions with stronger field.  Note that dependencies of the $ \\alpha $-effect, the turbulent magnetic diffusion coefficient and the effective drift velocity on the mean magnetic field for an isotropic turbulence have been found in \\cite{K91,RK93,KPR94} using a modified second order correlation approximation. Our results are different from that obtained in \\cite{K91,RK93,KPR94}. The reason is that in \\cite{K91,RK93,KPR94} a phenomenological procedure was used. In particular, in the first step of the calculations the nonlinear terms in the magnetohydrodynamic equations were dropped (which is valid for small hydrodynamic and magnetic Reynolds numbers or in a highly conductivity limit and small Strouhle numbers). In the next step of the calculations in \\cite{K91,RK93,KPR94} it was assumed that $ \\nu = \\eta = l_{c}^{2} / \\tau_{c} ,$ where $ l_{c} $ and $ \\tau_{c} $ are the correlation length and time of turbulent velocity field. The latter is valid when the hydrodynamic and magnetic Reynolds numbers are of the order of unit. In the present paper we use a different procedure (the $ \\tau $-approximation) for large hydrodynamic and magnetic Reynolds numbers.  In this study we also demonstrated an important role of two types of nonlinearities (algebraic and dynamic) in the mean-field dynamo. The algebraic nonlinearity is determined by a nonlinear dependence of the mean electromotive force on the mean magnetic field. The dynamic nonlinearity is determined by a differential equation for the magnetic part of the $ \\alpha $-effect. This equation is a consequence of the conservation of the total magnetic helicity (which includes both, the magnetic helicity of the mean magnetic field and the magnetic helicity of small-scale magnetic fluctuations). We found that at least for the $ \\alpha \\Omega $ axisymmetric dynamo the algebraic nonlinearity alone [{\\em i.e.,} the nonlinear functions $ \\alpha(B) ,$ $ \\eta_{A}(B) ,$ $ \\eta_{B}(B) $ and $ V_{A}(B) ]$ cannot saturate the dynamo generated mean magnetic field. The important parameter which characterizes the algebraic nonlinearity is the nonlinear dynamo number $ D(B) .$ The saturation of the growth of the dynamo generated mean magnetic field by the algebraic nonlinearity alone is possible when the derivative $ d D(B) / dB < 0 .$ We found that for the $ \\alpha \\Omega $ axisymmetric dynamo the nonlinear dynamo number $ D(B) $ is either a constant or $ D(B) \\propto B $ for $ B > B_{\\rm eq} / 3 $ depending on the model of the background turbulence. Therefore, in this case the algebraic nonlinearity alone cannot saturate the dynamo generated mean magnetic field.  The situation is changed when the dynamic nonlinearity is taken into account. The crucial point is that the dynamic equation for the magnetic part of the $ \\alpha $-effect ({\\em i.e.,} the dynamic nonlinearity) includes the flux of the magnetic helicity. Without the flux, the total magnetic helicity is conserved locally and the level of the saturated mean magnetic field is very low \\cite{KMRS2000}. The flux of the magnetic helicity results in that the total magnetic helicity is not conserved locally because the magnetic helicity of small-scale magnetic fluctuations is redistributed by a helicity flux. In this case an integral of the total magnetic helicity over the disc is conserved. The equilibrium state is given by a balance between magnetic helicity production and magnetic helicity transport \\cite{KMRS2000}. These two types of the nonlinearities (algebraic and dynamic) results in the equilibrium strength of the mean magnetic field is of order that of the equipartition field $ B_{\\rm eq} $ (see Section V) in agreement with observations of the galactic magnetic fields (see, {\\em e.g.,} \\cite{RSS88})."
    },
    "0206/astro-ph0206507_arXiv.txt": {
        "abstract": "{\\abstext%  \\abstext ",
        "introduction": "In recent years, cosmology has seen increasing observational evidence for an accelerating phase of the cosmic expansion, most notably through the observations of distant type Ia supernovae (Perlmutter et al.~1999, Riess et al.~1998). This astonishing evidence motivated renewed interest in the properties of the energy density ascribed to the ``vacuum''. A vacuum energy component should account for both the accelerating expansion and for the residual $\\sim70\\%$ of the energy density required for reconciling the geometrical flatness required by Cosmic Microwave Background (CMB) observations (De Bernardis et al.~2002, Lee et al.~2001, Halverson et al.~2002) with the evidence of a low-density universe with $\\Omega_0\\sim0.3$ (Percival et al.~2001).  While one of the historical candidates for such an energy density is the cosmological constant, introduced as a simple geometrical term in Einstein's equations, it is well known that it leads to serious and unsolved theoretical problems. The exceedingly low value of the vacuum energy density today, compared to that allowed by the most plausible theories of the early stage of the Universe, motivated the introduction of a more general concept now widely known as ``dark energy''.  Preceding the evidence for cosmic acceleration, a generalisation of the cosmological constant by means of a scalar field, now known as ``quintessence'', was proposed (Wetterich 1988, Ratra \\& Peebles 1988). In this class of models, the dark energy is supposed to reside mostly in the potential energy of the field, which interacts only gravitationally with ordinary matter. The evolution is described by the ordinary Klein-Gordon equation. If the potential is flat enough, or if the motion of the field along its trajectory is sufficiently slow, a cosmological constant-like behaviour can be mimicked by the scaling of the energy density of the quintessence field.  For general forms of the scalar field potential, there exist attractor trajectories for the evolution of the background expectation value of the field. These trajectories are known as ``tracking'' (Steinhardt, Wang \\& Zlatev 1999) and ``scaling'' (Liddle \\& Scherrer 1999) solutions. They have been shown to alleviate, at least at a classical level, the fine-tuning required in the early Universe, when the typical energy scales were presumably comparable to the Planck scale, to generate a vanishing relic vacuum energy as it is observed today, 120 orders of magnitude smaller. However, these scenarios are not able to solve the coincidence problem, i.e. why we are living in the epoch in which dark energy and matter have roughly the same energy density. Despite some attempts at addressing this issue (Tocchini-Valentini \\& Amendola 2002, Chiba 2001, Armendariz-Picon, Mukhanov \\& Steinhardt 2001, Dodelson, Kaplinghat \\& Stewart 2000), it remains one of the greatest puzzles of modern cosmology.  For constraining the nature of the dark energy, an important step would be accomplished if parameters could be constrained which capture its most essential features. In particular, if the dark energy is modelled as a quintessence field in the tracking regime, the simplest description of its properties will require the use of only two parameters, i.e.~its present energy density and the ratio between pressure and energy density in its equation of state. Generally, this ratio depends on time, as implied by the evolution according to the equation of motion. However, it can be shown that, in most simple models of quintessence involving an inverse power-law potential, the effect of a time variation of the equation of state can be neglected at low redshifts, when the field has settled on its tracking trajectory. In this case, the equation of state is simply related to the power with which the potential depends on the field itself. This simplification allows constructing a scheme for describing the dark energy behaviour at redshifts where interesting cosmological effects arise, such as the effect on the magnitude-redshift relation of type Ia supernovae (Perlmutter 1999, Riess 1998) and the effect on gravitational lensing of distant galaxies and quasars (Futamase \\& Yoshida 2001).  Even though the dark energy dynamics has a geometrical effect on acoustic features of the CMB anisotropy (Baccigalupi et al.~2002; Doran, Lilley \\& Wetterich 2002; Corasaniti \\& Copeland 2002), it is now commonly accepted that the most interesting properties of a dark energy component have to be probed by looking at processes occurring at low redshifts when it starts dominating the cosmic expansion.  Thus, one of the most powerful probes of the quintessence field results from its effects on cosmic structure formation, most notably on the background cosmology and the evolution of individual collapsing overdensities. First, a dark energy component affects the matter density of the background in which haloes form. Second, the amplitude of matter perturbations is sensitive to the presence of a dynamical vacuum energy, through the normalisation of the matter power spectrum to the large, unprocessed, scales probed by the large-scale CMB anisotropies. Third, changes in the background matter density induced by a dark energy component can seriously affect characteristic properties of collapsing structures. As shown by {\\L}okas \\& Hoffmann (2002), a substantial quintessence component changes the characteristic density of a forming dark matter halo.  In this paper, we study how quintessence affects the concentration of dark-matter haloes, and resulting changes in their weak-lensing efficiency. Weak gravitational lensing provides a powerful tool for mapping the large-scale mass distribution (see Mellier 1999a and Bartelmann \\& Schneider 2001 for reviews), and the potential impact of dark energy on the weak lensing convergence power spectrum has already been recognised (see Mellier 1999b; Huterer 2002). We show in this paper that an energy density component with negative pressure affects weak lensing by dark-matter haloes not only through changes in the global properties of the Universe, but also by modifying their internal density concentration. The main idea is that dark energy affects structure growth and thus the time of halo formation. Since halo concentrations reflect the density of the Universe at their formation epoch, this affects halo mass distributions, and weak lensing provides methods for quantifying respective changes.  The paper is organised as follows. In Sect.~2, we describe the main effects of the dark energy on the cosmological growth factor, volume elements, and the normalisation of the dark-matter power spectrum. In Sect.~3 we compute the impact on halo concentration. In Sect.~4 we predict resulting effects on the weak-lensing aperture mass statistics, and Sect.~5 contains our conclusions. ",
        "conclusions": "We investigated the expected properties of dark-matter haloes in dark energy cosmologies. For our purposes, the essential features of such models are captured describing the dark energy as a density component with negative pressure $p_\\mathrm{Q}=w\\rho_\\mathrm{Q}$, where $w\\ge-1$ is a constant. The dark energy density $\\rho_\\mathrm{Q}$ is determined by its present value in units of the critical density, $\\Omega_\\mathrm{Q}$. In agreement with results from observations of the CMB, we focus on models which are spatially flat, $\\Omega_0+\\Omega_\\mathrm{Q}=1$, have a matter density parameter $\\Omega_0=0.3$, and a Hubble constant of $h=0.7$.  The modified background dynamics in dark-energy models has two immediate consequences. First, the growth factor is changed, which determines how structures grow linearly against the expanding background. In models with fixed parameters $\\Omega_0$ and $\\Omega_\\mathrm{Q}$, structures form earlier if $w$ is larger. Second, under equal circumstances, the cosmic volume shrinks as $w$ increases. Dark energy models thus ``interpolate'' between low-density, spatially flat models with cosmological constant and low-density, open models.  For our purposes, we can neglect the clustering of the quintessence field and assume that the dark-matter power spectrum is given by the common CDM spectrum. We take the shape parameter to be given by $\\Gamma=\\Omega_0h$ and normalise the spectrum such that the COBE-DMR measurements of CMB fluctuations on large angular scales are reproduced. This implies a third cosmological consequence. As $w$ increases, the gravitational potential of matter fluctuations evolves more rapidly along a given line of sight. Secondary anisotropies in the CMB caused by the integrated Sachs-Wolfe effect thus grow in amplitude. Keeping the total fluctuation amplitude fixed to the COBE-DMR data thus requires the amplitude of the primordial fluctuations to decrease. Expressing the power-spectrum normalisation by $\\sigma_8$, this implies that $\\sigma_8$ must decrease as $w$ increases. The decrease is gentle for $-1\\le w\\lesssim-0.6$ and steepens as $w$ increases further.  These findings have two counter-acting effects on the evolution of dark-matter haloes. First, haloes forming earlier are more concentrated because their core density reflects the density of the background universe at their formation time. Since structures form earlier in dark energy models as $w$ is increased, haloes are expected to become more concentrated as $w$ grows. Second, the decrease of $\\sigma_8$ with increasing $w$ has the opposite effect on the halo formation time and indirectly on halo concentration. However, cosmologically interesting dark-energy equations of state must yield cosmic acceleration today and require $-1\\le w\\le -0.6$. Within that range, the first effect is dominant, and the overall behaviour is monotonic.  Ideally, extensive, high-resolution numerical simulations would be necessary for quantifying the net result of these two effects. However, simple algorithms for calculating halo concentrations have already been derived from existing numerical simulations. We used them for our work, assuming that they are also valid with the modification of Friedmann's equation caused by the introduction of dark energy instead of a cosmological constant.  We used three different recipes for computing halo concentrations. Albeit differing in detail, they agree in concept. Haloes are assigned a formation epoch, essentially requiring that a certain fraction of the final halo mass has already collapsed into sufficiently massive progenitors. The characteristic density of the haloes is then taken to be proportional to the mean background density of the universe at the halo formation epoch. We showed that all three recipes lead to the result that haloes are expected to be increasingly more concentrated as $w$ grows in quintessence models, showing that the effect of their earlier growth is stronger than the effect of decreasing $\\sigma_8$. This holds for $w\\lesssim-0.6$ and reverses for larger $w$ because the integrated Sachs-Wolfe effect then requires a steep decrease in $\\sigma_8$. The particular recipe for computing halo concentrations described by Bullock et al.~implies that haloes should be $\\sim50\\%$ more concentrated for $w=-0.6$ than for $w=-1$, where the increase is roughly linear with $w$.  Finally, we described that halo searches using weak-lensing techniques are sensitive to halo concentrations. Using the Sheth-Tormen modification of the Press-Schechter mass function for quantifying the halo population in mass and redshift, and the aperture mass technique for quantifying the weak-lensing effects of haloes, we showed that the expected number density on the sky of haloes causing 5-$\\sigma$ weak-lensing detections approximately doubles as $w$ increases from $-1$ to $-0.6$, where the increase is linear with $w$. Our results indicate that halo concentrations may be a sensitive probe for the dark-energy equation of state, and that gravitational lensing may provide the observational tools for applying that probe.  Note, however, that we did not allow variations in some cosmological parameters which may also change the number of weak-lensing haloes. In particular, an effect may arise from varying the index $n$ of the dark-matter power spectrum because it directly affects the determination of $\\sigma_8$. On the other hand, our approach here is to characterise the main effects of dark energy on halo formation and to propose weak lensing studies as a tool for constraining the dark energy itself, assuming the main cosmological parameters will be measured by independent observations like those of the CMB.  Thus, weak lensing turns out to be a powerful tool not only for mapping the distribution of matter in the Universe, but also for probing fundamental dark-energy properties. Currently, weak lensing observations do not allow detailed reconstructions of halo density profiles (Mellier 2001, Clowe et al.~2000; Mellier \\& Van Waerbeke 2001), mainly because of the resolution limit due to the finite number density of background galaxies.  On the other hand, interesting new perspectives have been opened by several recent wide-field cosmic-shear studies (Bacon, Refregier \\& Ellis 2000; Bacon, Massey, Refregier, Ellis 2002; Wittman et al.~2000, Van Waerbeke et al.~2000, Kaiser et al.~2000). Future weak lensing surveys will cover even larger fields and, thanks to the improved control of systematic errors, they will allow tighter constraints on the cosmological parameters. Besides the wide field telescopes which are currently in their project study phases, we specifically mention the ``dark matter telescope'' LSST (http://dmtelescop.org, proposed to scan a 7 square-degree sky field), the VISTA survey (http://www.vista.ac.uk), and the SNAP satellite (http://snap.lbl.gov), whose weak lensing survey will cover an area of 300 square degrees, resulting in a very wide field survey with excellent image quality and depth."
    },
    "0206/astro-ph0206488_arXiv.txt": {
        "abstract": "As shown by Parker and Raval, quantum field theory in curved spacetime gives a possible mechanism for explaining the observed recent acceleration of the universe. This mechanism, which differs in its dynamics from quintessence models, causes the universe to make a transition to an accelerating expansion in which the scalar curvature, $R$, of spacetime remains constant. This transition occurs despite the fact that we set the renormalized cosmological constant to zero. We show that this model agrees very well with the current observed type-Ia supernova (SNe-Ia) data. There are no free parameters in this fit, as the relevant observables are determined independently by means of the current cosmic microwave background radiation (CMBR) data. We also give the predicted curves for number count tests and for the ratio, $w(z)$, of the dark energy pressure to its density, as well as for $dw(z)/dz$ versus $w(z)$. These curves differ significantly from those obtained from a cosmological constant, and will be tested by planned future observations. ",
        "introduction": "\\label{sec:intro}  Observational evidence appears increasingly strong that the expansion of the universe is undergoing acceleration that started at a redshift $z$ of order 1 \\citep{Riess98,Riess01,Perlmutter98,Perlmutter99}. Observations of scores of type-Ia supernovae (SNe-Ia) out to $z$ of about 1.7 support this view~\\citep{Riess01}, and even glimpse the earlier decelerating stage of the expansion. It is fair to say that one of the most important questions in physics is: what causes this acceleration?  One of the more obvious possible answers is that we are observing the effects of a small positive cosmological constant $\\Lambda$ \\citep{Krauss95,Ostriker95,Dodelson96,Colberg00}. Another, less obvious possibility is that there is a quintessence field responsible for the acceleration of the universe \\citep{Caldwell98,Zlatev99,Dodelson00,Picon01}. Quintessence fields are scalar fields with potential energy functions that produce an acceleration of the universe when the gravitational and classical scalar (quintessence) field equations are solved. More recently, \\citet{Parker99a,Parker99b,Parker99c,Parker00,Parker01} showed that a quantized free scalar field of very small mass in its vacuum state may accelerate the universe. Their model differs from any quintessence model in that the scalar field is free, thus interacting only with the gravitational field. The nontrivial dynamics of this model arises from well-defined finite quantum corrections to the action that appear only in curved spacetime. This was the first model to present a realization of dark energy with ratio of pressure to energy density taking values more negative than $-1$. Other models having this property have subsequently been proposed \\citep{Caldwell02,Melchiorri02}.  The physics of the Parker-Raval model is based on quantum field theory in curved spacetime. The renormalized (i.e., observed) cosmological constant $\\Lambda$ is set to zero. Several mechanisms have been proposed that tend to drive the value of $\\Lambda$ to zero \\citep{Dolgov83,Ford87,Ford02,Tsamis98a,Tsamis98b,Abramo99}, but these mechanisms play no role in this model. The energy-momentum tensor of the quantized field in its vacuum state is determined by calculating an effective action \\citep{Schwinger51,DeWitt65,Jackiw74} in a general curved spacetime. The spacetime is unquantized, and is itself determined self-consistently from the Einstein gravitational field equations involving the vacuum expectation value of the energy-momentum tensor, as well as the classical energy momentum tensor of matter (including cold dark matter) and radiation. In solving the Einstein equations, the symmetries of the FRW spacetime are imposed, but the spacetime is not otherwise taken as fixed, and the initial value constraints of general relativity are satisfied. The acceleration is the result of including in the effective action a non-perturbative term involving the scalar curvature of the spacetime \\citep{Parker85a,Parker85b,Parker85c,Jack85}. The minimal effective action that includes this non-perturbative effect and gives the correct trace anomaly of the energy-momentum tensor was used by Parker and Raval. In applying this effective action to the recent expansion of the universe, terms involving more than two derivatives of the metric were neglected. In this approximation, a solution was proposed in which the universe undergoes a rapid transition from a standard FRW universe dominated by cold dark matter to one containing significant contributions of vacuum energy and pressure. The proposal is that this negative vacuum pressure is responsible for the observed acceleration of the universe. The reaction back of this negative vacuum pressure on the expansion of the universe is such as to cause a rapid transition to an expansion of the universe in which the scalar curvature remains constant. The vacuum pressure and energy density are determined by a single parameter related to the mass of the scalar field. The transition to constant scalar curvature is the result of a rapid growth in the magnitudes of the vacuum pressure and energy density that occurs in this theory when the scalar curvature approaches a particular value, of the order of the square of the mass of the particle associated with the scalar field. The Einstein equations cause a reaction back on the metric such as to prevent further increase in the magnitudes of the vacuum pressure and energy density. (This effect is analogous to Lenz's law in electromagnetism.)  The essential cosmological features of this model may be described quite simply. For times earlier than a time $t_j$ (corresponding to $z \\sim 1$), the universe undergoes the stages of the standard model, including early inflation, and radiation domination followed by domination by cold dark matter. During the latter stage, at time $t_j$ the vacuum energy and (negative) pressure of the free scalar quantized field increase rapidly in magnitude (from a cosmological point of view). The effect of this vacuum energy and pressure is to cause the scalar curvature $R$ of the spacetime to become constant at a value $R_j$. The spacetime line element is that of an FRW universe: \\begin{equation} ds^2 = -dt^2 + a(t)^2 [(1 - k r^2)^{-1} dr^2 + r^2 d\\theta^2 + r^2 \\sin^{2}{\\theta} d\\phi^2]\\;, \\label{ds2} \\end{equation} where $k= {\\pm 1}$ or $0$ indicates the spatial curvature. By joining at $t_j$ the matter dominated scale parameter $a(t)$ and its first and second derivatives to the solution for $a(t)$ in a constant $R$ universe, one uniquely determines the scale parameter $a(t)$ for times after $t_j$. This model is known as the vacuum cold dark matter (VCDM) model of Parker and Raval, or as the {\\em vacuum metamorphosis model} \\citep{Parker99c} to emphasize the existence of a rapid transition in the vacuum energy density and pressure.  The constant value, $R_j$, of the scalar curvature is a function of a single new parameter, $\\mbar$, related to the mass of the free scalar field. Therefore, the function $a(t)$ for $t > t_j$ is fully determined by $\\mbar$. The values of $\\mbar$ and $t_j$ can be expressed in terms of observables, namely, the present Hubble constant $H_0$, the densities $\\Omega_{m0}$ and $\\Omega_{r0}$ of the matter and radiation, respectively, relative to the closure density, and the curvature parameter $\\Omega_{k0} \\equiv -k/(H_0^2 a_0^2)$. (Here $a_0 \\equiv a(t_0)$ is the present value of the cosmological scale parameter.) These observables have been determined with reasonably good precision by various measurements that are independent of the SNe-Ia \\citep{Krauss00,Freedman01,Hu01,Huterer01,Turner01,Wang02}. Therefore, the value of $\\mbar$ is known to within narrow bounds, independently of the SNe-Ia observations.  The power spectrum of the CMBR depends largely on physical processes occurring long before $t_j$. The behavior of $a(t)$ in the VCDM model does not significantly differ from that of the standard model until after $t_j$. Therefore, the predicted power spectrum in the VCDM model differs only slightly from that of the standard model. We calculate the predicted power spectrum of the CMBR in the VCDM model (as described below), and find the range of values of the above observables that give a good fit to the CMBR observations.  From this range of observables, the corresponding range of the parameter $\\mbar$ follows. Therefore, the prediction of the VCDM model for the magnitude versus redshift curve of the SNe-Ia is completely determined, with no adjustable parameters. We plot the predicted curves (obtained from this range of $\\mbar$) for the distance modulus $\\Delta(m - M)$ of the SNe-Ia as a function of $z$. Comparison with the observed data points, as summarized by \\citet{Riess01}, shows that a significant subset of predicted curves fit the SNe-Ia data very well, passing within the narrow error bars of each of the binned data points, as well as of the single data point at $z \\approx 1.7$.  We also give the curves predicted by the VCDM model for number counts of cosmological objects as a function of $z$ and for the ratio, $w(z)$, of the vacuum pressure to vacuum energy density, as well as for $dw/dz$ versus $w$ (parametrized by $z$). The predictions of the VCDM model differ significantly from those of the $\\Lambda$CDM model. Accurate measurements of these quantities out to $z$ of about 2 would be very telling.  The model we consider here is the simplest of a class of models in which a transition occurs around a finite value of $z$ within the range of possible observation. This general class has been studied from a phenomenological point of view by \\citet{Bassett02a,Bassett02b}. They find that the CMBR, large scale structure, and supernova data tend to favor a late-time transition over the standard $\\Lambda$CDM model. Although they considered only dark-energy equations of state, $w$, with $w > -1$, their phenomenological analysis can be generalized to include the present VCDM model. In addition, the VCDM model is readily generalized to include a nonzero vacuum expectation value of the low mass scalar field, which could bring $w$ into the range greater than $-1$. In the present paper, we are taking the simplest of the possible VCDM (or vacuum metamorphosis) models, so as to introduce no arbitrary parameters into the fit to the supernova data. ",
        "conclusions": "We have shown that the current observational data indicating that the expansion of the Universe is undergoing acceleration are quite consistent with the hypothesis that a transition to a constant-scalar-curvature stage of the expansion occurred at a redshift $z\\sim 1$ in the spatially flat FRW universe having zero cosmological constant. This is the scenario proposed in the VCDM (or vacuum metamorphosis) model introduced by Parker and Raval. The late constancy of the scalar curvature at a value $R_j= \\mbar^2$ is induced by quantum effects of a free scalar field of low mass in the curved cosmological background. The parameter $\\mbar$, related to the mass of the field, is the only new relevant parameter introduced in this model, and can be expressed in terms of the present cosmological parameters $H_0$, $\\Omega_{m0}$, $\\Omega_{r0}$, and $\\Omega_{k0}$ (see eq.~[\\ref{mHdimensionless}]).  Comparison of the CMBR-power-spectrum data with the flat-VCDM-model prediction, without or with the HST-Key-Project result as a constraint (see figs.~\\ref{fig:CMBR} and \\ref{fig:95cl}, and table~\\ref{tab:peaks}), gives the values of the cosmological parameters to be $H_0=65^{-(16;1)}_{+10}$~${\\rm km}$~${\\rm s}^{-1}$~${\\rm Mpc}^{-1}$ and $\\Omega_{m0}=0.34^{+(0.46;0.08)}_{-0.14}$. (Recall the definition of our notation in sec.~\\ref{sec:mbar}: the uncertainties appearing in parenthesis refer to the $95\\%$ confidence level, without and with the HST constraint, respectively.) Such values lead to $\\mbar= 4.52^{-(1.76;0.18)}_{+0.84} \\times 10^{-33}$~${\\rm eV}$, and the best-fit values from the CMBR data give rise to a very good no-parameter fit to the SNe-Ia observational data. However, the SNe-Ia data are not accurate enough to draw a clear distinction between the VCDM and $\\Lambda$CDM models. Other quantities of interest predicted by the VCDM model with the cosmological parameters mentioned above are the time and redshift at the transition between the matter-dominated and constant-scalar-curvature stages, ($t_j=5.33^{+(3.41;0.22)}_{-0.84}$~${\\rm Gyr}$ and $z_j=1.19^{-(0.76;0.19)}_{+0.47}$), the time and redshift when the accelerated expansion started ($t_a=8.0^{+(5.2;0.4)}_{-1.2}$~${\\rm Gyr}$ and $z_a=0.67^{-(0.59;0.15)}_{+0.35}$), and the age of the Universe, $t_0=14.9\\mp 0.8$~${\\rm Gyr}$.  Regarding future tests of the VCDM model, we have presented the prediction of number counts as a function of redshift, and compared it with the analogous $\\Lambda$CDM prediction (see fig.~\\ref{fig:dndz}). For approximately the same cosmological parameters, the VCDM model predicts nearly $30\\%$ more objects to be observed in a small redshift interval around $z\\approx 1$ than the $\\Lambda$CDM model. Data provided by the DEEP Redshift Survey in the near future will likely be able to distinguish these two models. Also, DEEP data combined with future measurements of SNe-Ia luminosity distances provided by the proposed SNAP satellite should greatly improve our knowledge of the dark energy equation of state, which bears the most distinct feature of the VCDM model: $w<-1$ and $w'<0$ (see figs.~\\ref{fig:wz} and \\ref{fig:dwdz}).  It should be noted that we have here considered the simplest form of the VCDM model, in which the transition to constant scalar curvature is continuous and effectively instantaneous (see fig.~\\ref{fig:wT}). This form of the model makes definite predictions regarding the distance moduli of SNe-Ia and number counts. Thus, it is encouraging that it remains a viable model when confronted with the current observational data. Other natural parameters that may come into the VCDM model are the time interval over which the transition occurs, and the vacuum expectation value of the scalar field. A nonzero value of the transition time interval would mainly affect the predictions around $z\\sim 1$, and a nonzero value of the vacuum expectation value is likely to increase the ratio of pressure to density, $w$. Future observational data will determine if it is necessary to consider nonzero values for these parameters."
    },
    "0206/astro-ph0206441_arXiv.txt": {
        "abstract": "Recent discoveries by the Sloan Digital Sky Survey (SDSS) of four bright $z \\sim 6$ quasars could constrain the mechanism by which the supermassive black holes powering these sources are assembled. Here we compute the probability that the fluxes of the quasars are strongly amplified by gravitational lensing, and therefore the likelihood that the black hole masses are overestimated when they are inferred assuming Eddington luminosities. The poorly--constrained shape of the intrinsic quasar luminosity function (LF) at redshift $\\sim 6$ results in a large range of possible lensing probabilities.  If the LF is either steep, or extends to faint magnitudes, the probability for amplification by a factor $\\mu\\ga10$ (and with only one image detectable by SDSS) can reach essentially $100 \\%$.  We show that future observations, in particular, either of the current four quasars at the high angular resolution provided by the Hubble Space Telescope, or of an increased sample of $\\sim 20$ redshift six quasars at the current angular resolution, should either discover several gravitational lenses, or else provide interesting new constraints on the shape of the $z\\sim 6$ quasar LF. ",
        "introduction": "\\label{sec:intro} The Sloan Digital Sky Survey (SDSS) has recently discovered several extremely bright and distant quasars with redshifts $z \\sim 6$, whose luminosities are a result of gas accretion onto supermassive black holes (BHs). These rare objects have luminosities at the far bright tail of the quasar luminosity function (LF), corresponding to $\\sim 5 \\; \\sigma$ peaks in the density field in models in which quasars populate dark matter halos (Haiman \\& Loeb 2001; Fan et al. 2001b). By assuming that the observed luminosities equal the limiting Eddington luminosities of the BHs, the masses of all these BHs are inferred to be a few $\\times 10^9~{\\rm M_\\odot}$.  The mere existence of such massive BHs at so early a stage in the evolution of the universe should provide insight into the formation and growth of supermassive black holes (Haiman \\& Loeb 2001).  The above estimates for the masses of the BHs associated with the SDSS quasars scale directly with the inferred luminosities, and thus with the observed fluxes of the quasars. However, the apparent fluxes can be increased by gravitational lensing. If a quasar is gravitationally lensed and this is not taken into account in the mass estimate, the quasar's intrinsic luminosity, and therefore its BH mass, will be overestimated. Previous considerations of the probability that these SDSS quasars are lensed indicate that while the a~priori lensing probability is small, the magnification bias is poorly known and can be quite large (Wyithe \\& Loeb 2002).  If massive BHs grow from stellar--mass seeds by accreting gas at the Eddington rate, they reach masses of $\\sim 10^9~{\\rm M_\\odot}$ in $\\sim 10^9$ years, a time span comparable to the age of the universe at $z \\sim 6$ in the current ``best--fit'' cosmologies.  Since the presence of such massive BHs can only be marginally accommodated into structure formation models (Haiman \\& Loeb 2001), it is interesting to ask whether {\\it{all}} of the SDSS quasars may be strongly lensed.  In this paper, we compute the expected probabilities for lensing of $z\\sim6$ quasars by intervening galaxies. We quantify how the probability depends on the shape of the assumed quasar LF, and in particular, we focus on the question: can the expected probabilities reach values near unity? We show that current constraints on the $z\\sim 6$ quasar LF are sufficiently weak that the probabilities for strong lensing can indeed be close to $100\\%$. We will utilize this result to ``invert`` the problem, and demonstrate how the intrinsic LF can be constrained using current and future detections (or absence) of lensing events with multiple images.  The rest of this paper is organized as follows. In \\S~2, we describe our model for the population of gravitational lenses. In \\S~3, we derive the resulting intrinsic ``a~priori'' lensing probabilities along a random line of sight to redshift $z\\sim 6$.  In \\S~4, we discuss the effects of magnification bias, and the expected a~posteriori lensing probabilities for the $z\\sim 6$ quasars in models with different LFs.  In \\S~5, we quantify constraints on the quasar LF that can be derived from current and future searches for lensing events. Finally, in \\S~6, we discuss our main results and then summarize the implications of this work.  Throughout this paper, we adopt a flat cosmological model dominated by cold dark matter (CDM) and a cosmological constant ($\\Lambda$), with $\\Omega_m=0.25$, $\\Omega_b=0.04$, and $\\Omega_\\Lambda=0.75$, a Hubble constant $H_0=70~{\\rm km~s^{-1}}$, an rms mass fluctuation within a sphere of radius $8 \\; h^{-1}$ Mpc of $\\sigma_8=0.9$, and power--law index $n=1$ for the power spectrum of density fluctuations. We also adopt the cosmological transfer function from Eisenstein \\& Hu (1999). ",
        "conclusions": "In this paper, we computed the expected lensing probabilities of $z\\sim6$ quasars by intervening galaxies. The a~posteriori probability that a $z \\sim 6$ quasar is strongly lensed (by a factor $\\mu>10$ in flux enhancement), but without producing two images detectable in the Sloan Survey, can be very large, depending on the shape of the LF. The LFs we use are consistent with other available observational constraints, yet these constraints still allow for a wide range of possibilities.  Though an {\\it observed} bright-end slope of $\\beta_h=4.3$ could be ruled out at the 99 percent confidence level (Fan et al.\\ (2001)), lensing alters the slope so that the apparent LF will have a slope close to $\\sim 3.0$ if magnification bias is important. As illustrated in Fig.~\\ref{fig:lfobs}, even the steepest intrinsic LF we consider, with a bright end slope of $4.5$, is consistent with the SDSS limit on slope. This will be true for arbitrarily steep slopes, provided that the characteristic luminosity is not increased.  In a similar probability calculation, Wyithe \\& Loeb (2002) concluded that the observed probability reaches $30 \\%$ that a single $z \\sim 6$ quasar is lensed and magnified by a factor of 10 or more. If we restrict our analysis to the choices of LFs considered in their paper, we agree with this result. However, we here use a wider range of quasar LFs and conclude that the lensing probability can reach essentially 100\\%.  As a result, the current observations are consistent with all four $z \\sim 6$ SDSS quasars being strongly lensed.  If the probability that all four quasars are lensed were high, then this would alleviate the problematic time constraints on assembling supermassive black holes at the earliest stages in the evolution of the universe . However, in a separate analysis, based on the apparently large size of the ionized region around the SDSS quasar at $z=6.28$, Haiman \\& Cen (2002) place a strong constraint on the lensing magnification of this source, and find $\\mu<5$. This result depends primarily on the assumption that the source is embedded in a neutral (rather than ionized) intergalactic medium.  We have also considered the probabilities for lensing events that would have been detectable by SDSS. The lack of such detections can be used to place constraints on the quasar LF.  Although constraints from the current four quasars are mild, the situation is likely to improve.  As illustrated in Fig.~\\ref{fig:contour0}, increasing the SDSS sample from 4 to 20 objects, something that is likely to happen over the next few years, would allow one to rule out interesting quasar LF models. The expected probability of detecting multiple images is also sensitive to the angular resolution of the observations. In terms of the constraints on the shape of the intrinsic LF, observing four quasars at a resolution of $0.1 ''$ is roughly equivalent to observing 20 objects at the resolution of the SDSS ($1 ''$).  Hence, upcoming observations, in particular with the \\emph{Hubble Space Telescope}, should reveal whether a significant fraction of the $z\\sim 6$ quasars is lensed, and will allow us to place strong constraints on their intrinsic LF."
    },
    "0206/astro-ph0206394_arXiv.txt": {
        "abstract": "We have performed 3-D numerical magnetohydrodynamic (MHD) jet experiments to study the instabilities associated with strongly toroidal magnetic fields. In contemporary jet theories, a highly wound up magnetic field is a crucial ingredient for collimation of the flow. If such magnetic configurations are as unstable as found in the laboratory and by analytical estimates, our understanding of MHD jet driving and collimation has to be revised.  A perfectly conducting Keplerian disc with fixed density, rotational velocity and pressure is used as a lower boundary for the jet. Initially, the corona above the disc is at rest, permeated by a uniform magnetic field, and is in hydrostatic equilibrium in a softened gravitational field from a point mass. The mass ejection from the disc is subsequently allowed to evolve according to deviations from the initial pressure equilibrium between disc and corona.  The energy equation is solved, with the inclusion of self-consistently computed heating by viscous and magnetic dissipation. We find that magnetic dissipation may have profound effects on the jet flow as: 1) it turns on in highly wound up magnetic field regions and helps to prevent critical kink situations; 2) it influences jet dynamics by re-organizing the magnetic field structure and increasing thermal pressure in the jet; and 3) it influences mass loading by increasing temperature and pressure at the base of the jet.  The resulting jets evolve into time-dependent, non-axisymmetric configurations, but we find only minor disruption of the jets by e.g.\\ the kink instability. ",
        "introduction": "Though astrophysical jets have been observed in a rather wide range of accreting systems, it is generally assumed that the mechanism for acceleration and collimation is generic \\cite[for recent reviews]{Livio,spruit00}. Contemporary jet theories rely on magnetic forces as the jet producing mechanism, either with the flow emanating from the disc surface and associated with an open magnetic field structure \\cite{BP82} or with the flow emanating from the disc-star interface and associated with a closed field structure connecting disc and central star \\cite{pringle}. In the former so-called disc-wind scenario, which is currently receiving most attention, the driving force, acting to overcome the gravitational pull of the central object, is typically regarded as a component of the centrifugal force along the magnetic field. However, adopting an inertial frame of reference the driving mechanism may equally be interpreted as magnetic \\cite{spruit96}. The inertia of the outflowing gas eventually becomes dynamically important and the gas stops corotating with the underlying disc. At this point the magnetic field, which cannot easily slip through the highly ionized gas, is wound up in a cork screw manner between the disc and the vertically accelerated and less quickly rotating outflow. The tension (``hoop stress'') of the wound-up magnetic field configuration is generally believed to be the collimating force, but collimation by a poloidal field has been proposed as well \\cite{Spruit} motivated by instability arguments against the wound-up field scenario.  The assumed large scale magnetic field in the disc-wind scenario may either be provided by dynamo processes in the disc \\cite{Brandenburg2000,Brandenburg+2000} or captured from the environment and advected inwards by the accreting matter \\cite{CaoSpruit}. Numerical work has been carried out to investigate situations in which a significant fraction of the field lines loop back into the disc \\cite{Romanovab,TBR}. These experiments have focused on the ejection and acceleration mechanisms as well as the evolution of the magnetosphere close to the disc, and have not followed the flow further out for extended times.  The numerical work done in relation to the ``classical'' large scale open field scenario may be divided into experiments attempting to include (parts of) the accretion disc in the computational domain \\cite{US85,BL95} and more recent work \\cite{Ustyugova,OPS,Meier} which has used the disc only as a fixed lower boundary to avoid problems with radial collapse of the disc and thereby relatively short time evolution of the experiment. Both types of experiments have been axisymmetric and as such have not questioned the potentially (kink) unstable wound up magnetic field configuration on which the collimation process relies. However, Spruit et al.\\ \\cite*{spruit+00} have recently proposed fast reconnection processes in more disordered non-axisymmetric large scale magnetic fields in the jet like gamma ray burst (GRB) scenario. The magnetic dissipation is proposed for the GRB fireballs e.g.\\ to produce the observed radiation with better efficiency. We find that such magnetic processes may have severe impact on jet dynamics and stability in particular.  The main purpose of the work presented here is to establish whether or not the jets in the disc-wind scenario are prone to catastrophic MHD instabilities. This calls for an implementation of a setup in three dimensions which will be described in Sect.~\\ref{sec:model}. For simplicity, we assume a large-scale open magnetic field structure and use the disc as a fixed boundary. The details of the initial conditions and boundary conditions are found in Sect.~\\ref{sec:ic} and \\ref{sec:bc} respectively. Results are presented in Sect.~\\ref{sec:results}, with Sect.~\\ref{sec:dynamics} concentrating on the jet dynamics and Sect.~\\ref{sec:stability} on the observed 3-D jet stability in the experiments. In Sect.~\\ref{sec:discussion} we discuss the results with special emphasis on thermal properties and magnetic field structure. Finally, the paper is summarized and conclusions are presented in Sect.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  A high order numerical scheme has successfully been adopted and a suitable mesh refinement specified. Initial conditions resembling previous axisymmetric numerical experiments have been chosen to ease comparison. The polytropic equation of state is only used initially, to prescribe the initial density distribution of the corona. The most important features that differ from previous jet experiments are: \\begin{itemize} \\item The model is three dimensional rather than axisymmetric. Due to the periodic boundaries, a cutoff of the disc at large radii is applied. \\item The thermal energy equation is solved, with self-consistently computed heating by viscous and Joule dissipation. \\item ``Free'' mass outflow from the disc, i.e.\\ the mass flux is allowed to adjust self-consistently to the forces near the disc surface. \\item Parameterized Poynting flux through the upper boundary, representing external work. \\item The experiments have been evolved far beyond the initial transient, and display a quasistationary behavior, as evidenced for example by the nearly constant total energy flux. \\end{itemize}  The jet dynamics has been investigated by analyzing the mechanisms of energy conversion. The rotational energy of the disc is carried by the magnetic field into the corona and is first predominantly converted into kinetic energy. In the upper two thirds of the computational domain the magnetic energy is predominantly converted into thermal energy. General features predicted by steady state theory and axisymmetric numerical experiments, such as knot generation and terminal velocity dependency on the magnetic field strength have been confirmed in the relatively early and quiescent stages of the experiments. At later stages the flow becomes quite unsteady as instabilities set in, but no serious disruption of the flow occurs. The jet stability is found to be influenced by the magnetic dissipation --- this has not previously been investigated in the context of jet flows. The main findings concerning magnetic dissipation are the following: \\begin{itemize} \\item The heating by magnetic dissipation is significant, and leads to jet temperatures of the order of the virial temperature of the innermost Kepler orbit. \\item Magnetic reconnection occurs primarily in two regions: in a central region of the jet due to field interlocking and in a jet cocoon due to the spiraling motion of the jet beam forcing the wound up field into the ambient vertical field. \\item Magnetic dissipation helps to prevent critical kinking, and the jet is able to sustain a quasistationary flow, with only limited excursions. \\item Reconnection events are seen to result in mass ejection into the ambient medium and cause filament structures in the jet beam. \\item Reconnection events are seen to significantly influence the dynamics at the jet flanks where deceleration and even back-flow can be found quite close to the central super-Alfv\\'{e}nic beam. \\item Heating in the region just above the disc is likely to have a significant effect on the mass loading. \\end{itemize}"
    },
    "0206/astro-ph0206306_arXiv.txt": {
        "abstract": "{We combine the catalog of compact high--velocity \\hi clouds extracted by de\\,Heij et al. (\\cite{deheij02}) from the Leiden/Dwingeloo Survey in the northern hemisphere with the catalog extracted by Putman et al. (\\cite{putman02a}) from the Parkes HIPASS data in the southern hemisphere, and analyze the all--sky properties of the ensemble. Compact high--velocity clouds are a subclass of the general high--velocity cloud phenomenon which are isolated in position and velocity from the extended high--velocity Complexes and Streams down to column densities below 1.5$\\times10^{18}$\\,cm$^{-2}$. Objects satisfying these criteria for isolation are found to have a median angular size of less than one degree.  We discuss selection effects relevant to the two surveys; in particular the crucial role played by obscuration due to Galactic H\\,{\\sc i}.  Five principal observables are defined for the CHVC population: \\,(1) the spatial deployment of the objects on the sky, (2) the kinematic distribution, (3) the number distribution of observed \\hi column densities, (4) the number distribution of angular sizes, and (5) the number distribution of \\hi linewidth.  Two classes of models are considered to reproduce the observed properties.  The agreement of models with the data is judged by extracting these same observables from simulations, in a manner consistent with the sensitivities of the observations and explicitly taking account of Galactic obscuration.  We show that models in which the CHVCs are the \\hi counterparts of dark--matter halos evolving in the Local Group potential provide a good match to the observables. The best--fitting populations have a maximum HI mass of $10^7\\rm\\;M_\\odot$, a power-law slope of the HI mass distribution in the range $-1.7$ to~$-1.8$, and a Gaussian dispersion for their spatial distributions of between 150 and 200~kpc centered on both the Milky Way and M\\,31. Given its greater mean distance, only a small fraction of the M\\,31 sub--population is predicted to have been detected in present surveys. An empirical model for an extended Galactic halo distribution for the CHVCs is also considered. While reproducing some aspects of the population, this class of models does not account for some key systematic features of the population. ",
        "introduction": "Since the discovery of the \\hi high--velocity clouds by Muller et al. (\\cite{muller63}), different explanations, each with its own characteristic distance scale, have been proposed. It is likely that not all of the anomalous--velocity \\hi represents a single phenomenon, in a single physical state.  Determining the topology of the entire population of anomalous--velocity \\hi is not a simple matter, and the task is all the more daunting to carry out on an all--sky basis because of disparities between the observational survey material available from the northern and southern hemispheres.  The question of distance remains the most important, because the principal physical parameters depend on distance: mass varying as $d^2$, density as $d^{-1}$, and linear size directly as $d$. Most of the \\hi emission at anomalous velocities is contributed from extended complexes containing internal sub-structure but embedded in a common diffuse envelope, with angular sizes up to tens of degrees.  Such structures include the Magellanic Stream of debris from the Galaxy/LMC interaction and several HVC complexes, most notably complexes A, C, and H.  The complexes are few in number but dominate the \\hi flux observed.  The Magellanic Stream comprises gas stripped from the Large Magellanic Cloud, either by the Galactic tidal field or by the ram--pressure of the motion of the LMC through the gaseous halo of the Galaxy.  It therefore will be located at a distance comparable to that of the Magellanic Cloud, i.e. some $50\\rm\\;kpc$ (see e.g. Putman \\& Gibson \\cite{putm99}).  The distance to Complex A has been constrained by van Woerden et al. (\\cite{woer99}) and then more tightly by Wakker (\\cite{wakker01}) to lie within the distance range $8<d<10$ kpc.  If, as seems plausible, the other large complexes also lie at distances ranging from several to some 50 kpc, they will have been substantially affected by the radiation and gravitational fields of the Milky Way.  Another category of anomalous \\hi high--velocity clouds are the compact, isolated high--velocity clouds discussed by Braun \\& Burton (\\cite{braun99}).  CHVCs are distinct from the HVC complexes in that they are sharply bounded in angular extent at very low column density limits, i.e. below 1.5$\\times$10$^{18}$\\,cm$^{-2}$ (de\\,Heij et al. \\cite{deheij02}). This is an order of magnitude lower than the critical \\hi column density of about 2$\\times$10$^{19}$\\,cm$^{-2}$, where the ionized fraction is thought to increase dramatically due to the extragalactic radiation field. For this reason, these objects are likely to provide their own shielding to ionizing radiation. Although not selected on the basis of angular size, such sharply bounded objects are found to be rather compact, with a median angular size of less than 1~degree.  An analogy of the CHVC ensemble with that of the dwarf galaxy population in the Local Group is suggestive, and illustrates the hypothesis that is under discussion here.  Some few Local Group dwarf galaxies also extend over large angles.  The Sgr Dwarf Spheroidal discovered by Ibata et al.  (\\cite{ibat94}) spans some $40\\deg$; presumably it was once a rather conventional dwarf, but its current proximity to the Milky Way accounts for its large angular size. This proximity has fundamentally distorted its shape, and will determine its further evolution.  The streams of stars found in the halo of the Galaxy by Helmi et al. (\\cite{helmi99}) probably represent even more dramatic examples of the fate which awaits dwarf galaxies which transgress into the sphere of the Galaxy's dominance.  Analogous stellar streams have been found in M\\,31 by Ibata et al.  (\\cite{ibat01}) and by Choi et al. (\\cite{choi02}), as well as in association with Local Group dwarf galaxies by Majewski et al.  (\\cite{maje00}), indicating that accretion (and subsequent stripping) of satellites is an ongoing process.  If a selection were to be made of the dwarf galaxy population in the Local Group on the basis of angular size, the few large--angle systems which would be selected, namely the LMC and SMC, and the Sgr dwarf, and --- depending on the flexibility of the selection criteria --- perhaps the ill--fated coherent stellar streams in the Galactic halo, would represent systems nearby, of large angular extent, and currently undergoing substantial evolution.  Those systems selected on the basis of being compact and isolated, on the other hand, would represent dwarf galaxies typically at substantial distances and typically at a more primitive stage in their evolution. Regarding distance and evolutionary status they would differ from the nearby, extended objects, although at some earlier stage the distinction would not have been relevant.  The possibility that some of the high--velocity clouds might be essentially extragalactic has been considered in various contexts by, among others, Oort (\\cite{oort66}, \\cite{oort70}, \\cite{oort81}), Verschuur (\\cite{vers75}), Eichler (\\cite{eich76}), Einasto et al. (\\cite{eina76}), Giovanelli (\\cite{giov81}), Bajaja et al. (\\cite{baja87}), Wakker \\& van Woerden (\\cite{wakker97}), Braun \\& Burton (\\cite{braun99}), and Blitz et al. (\\cite{blitz99}).  It is interesting to note that the principal earlier arguments given against a Local Group deployment, most effectively in the papers cited above by Oort and Giovanelli, were based on the angular sizes of the few large complexes and on the predominance of negative velocities in the single hemisphere of the sky for which substantial observational data were then available. The more complete data available now, however, show about as many features at positive velocities as at negative ones.  It is also interesting to note that the papers by Eichler and by Einasto et al. cited above consider distant high--velocity clouds as possible sources of matter, including dark matter, fueling continuing evolution of the Galaxy.  Blitz et al. (\\cite{blitz99}) revived the suggestion that high--velocity clouds are the primordial building blocks fueling galactic growth and evolution, and argued that the extended complexes owe their angular extent to their proximity.  Braun \\& Burton (\\cite{braun99}) identified CHVCs as a subset of the anomalous--velocity gas that might be characteristic of a single class of HVCs, whose members plausibly originated under common circumstances and share a common subsequent evolutionary history. They emphasized the importance of extracting a homogenous sample of independently confirmed objects from well--sampled, high--sensitivity \\hi surveys.  The spatial and kinematic distributions of the CHVCs were found by Braun \\& Burton to be consistent with a dynamically cold ensemble spread throughout the Local Group, but with a net negative velocity with respect to the mean of the Local Group galaxies.  They suggested that the CHVCs might represent the low--circular--velocity dark matter halos predicted by Klypin et al.  (\\cite{klyp99}) and Moore et al. (\\cite{moor99}) in the context of the hierarchical structure paradigm of galactic evolution. These halos would contain no, or only a few, stars; most of their visible matter would be in the form of atomic hydrogen.  Although many of the halos would already have been accreted into the Galaxy or M\\,31, some would still populate the Local Group, either located in the far field or concentrated around the two dominate Local Group galaxies. Those passing close to either the Milky Way or M\\,31 would be ram--pressure stripped of their gas and tidally disrupted by the gravitational field. Near the Milky Way, the tidally distorted features would correspond to the high--velocity--cloud complexes observed.  The quality and quantity of survey material is important to interpretation of the CHVC population, which is a global one.  The observational data entering this analysis involved merging two catalogs of CHVCs, one based on the material in the Leiden/Dwingeloo Survey (LDS) of Hartmann \\& Burton (\\cite{hartmann97}), and the other based on the \\hi Parkes All--Sky Survey (HIPASS) described by Barnes et al. (\\cite{barnes}). Both of these surveys were searched for anomalous--velocity features using the algorithm described by de\\,Heij et al. (\\cite{deheij02}, Paper~I).  This algorithm led to the LDS catalog of de\\,Heij et al. for the CHVCs at declinations north of $-30\\deg$, and to the HIPASS catalog of Putman et al. (\\cite{putman02a}) for those at $\\delta<0\\deg$. The surveys overlap in the declination range $+2\\deg <\\delta < -30\\deg$, allowing estimates of the relative completeness of the catalogs.  We are able to predict how a survey with the LDS parameters would respond to the CHVCs detected by HIPASS, and vice versa. In the subsequent simulations, we are able to {\\emph {sample the simulated material as if it were being observed}} by one of these surveys.  This paper is organized as follows. We describe the application of the algorithm in \\S\\,\\ref{sec:algorithm}. In \\S\\,\\ref{sec:selection} we discuss various observational selection effects, and indicate how to account for these. We address obscuration by our own Galaxy in \\S\\,\\ref{sec:obscur}, the consequences of the differing observational parameters of the LDS and HIPASS data in \\S\\,\\ref{sec:obspara}, and the resulting differing degrees of completeness of the LDS and HIPASS catalogs in \\S\\,\\ref{sec:complete}.  We discuss the observable all--sky properties of the CHVC ensemble in \\S\\,\\ref{sec:allsky}, including the spatial deployment (in \\S\\,\\ref{sec:spatial}), the kinematic deployment (in \\S\\,\\ref{sec:velocity}), and the distributions of \\hi flux and angular size (in \\S\\,\\ref{sec:over}).  We then attempt to reproduce these properties by considering models, first based on Local Group distributions as discussed in \\S\\,\\ref{sec:model} and then on distributions within an extended Galactic Halo as discussed in \\S\\,\\ref{sec:simplemodel}; these simulations are sampled {\\emph {as if being observed}} with the LDS and HIPASS programs. We discuss the conclusions that can be drawn from this analysis in \\S\\,\\ref{sec:conclusion}.   \\begin{figure*} \\sidecaption \\caption{Fraction of a homogeneously distributed sample of test clouds that is not blended with Galactic emission.  The sample was distributed on the sky with a Gaussian velocity distribution in the Local Standard of Rest system (upper panel), the Galactic Standard of Rest system (middle), and  the Local Group Standard of Rest (lower). The average velocity of the Gaussian velocity distributions is $-50\\rm\\;km\\;s^{-1}$ for each of the simulations; the velocity dispersion is $240\\rm\\;km\\;s^{-1}$ for the LSR representation, and $110\\rm\\;km\\;s^{-1}$ for the GSR and LGSR ones, in rough accordance with the observed situation. The obscuration by Galactic emission was simulated by removing that part of the sample for which the deviation velocity (measured with respect to the LSR) is less than $70\\rm\\;km\\;s^{-1}$, based on a model of Galactic kinematics which incorporates observed properties of the gaseous disk, including its warp and flare.  The apparent structures are a consequence of the non--uniform obscuration in position and velocity. The appearance of the \\hi Zone of Avoidance differs when material is considered in different reference frames.} \\label{fig:calcskydistr} \\end{figure*}  \\begin{figure*} \\centering \\caption{Demonstration of the influence of Galactic obscuration on observable properties of a CHVC population. The left--hand panels show the distributions of measured (after obscuration) versus actual mean velocity and velocity dispersion determined from a series of 1000 simulations of a population of 200 objects.  The upper panel on the left represents a population with a Gaussian velocity distribution in the GSR reference frame, with $\\mu = -50\\rm\\;km\\;s^{-1}$ and $\\sigma = 115\\rm\\;km\\;s^{-1}$; the lower left--hand panel represents a population with a Gaussian velocity distribution in the LGSR frame, with $\\mu = -55\\rm\\;km\\;s^{-1}$ and $\\sigma = 105\\rm\\;km\\;s^{-1}$.  Obscuration removes about 30\\% of the population and leads to both an overestimate of the dispersion and a more negative estimate of the mean velocity. The right--hand panels illustrate the degree to which populations in the GSR and LGSR frames could be distinguished via their statistical parameters. Measured and actual parameter differences between the GSR and LGSR frames are contrasted for 1000 simulated populations of 200 objects, half defined in the GSR frame and half in the LGSR frame.  All populations have a mean velocity of~$-50\\rm\\;km\\;s^{-1}$ and a dispersion of~$110\\rm\\;km\\;s^{-1}$ in their reference frame.  The measured parameter differences for the observed CHVC sample are indicated by the dashed lines. While the mean velocity does not provide significant distinguishing capability between the GSR and LGSR frames, the velocity dispersion does. }\\label{fig:obscure} \\end{figure*}  \\begin{figure*} \\centering \\caption{~Demonstration of the effects of differing resolution and sensitivity in the LDS and HIPASS data on the extracted  CHVC samples.~~{\\it upper left:} \\,Comparison of the velocity FWHM for the CHVCs found in the HIPASS catalog (black line) with the velocity widths of those found in the LDS catalog  (red line).  The velocity resolution of the LDS is 1.03 \\kms, but 26 \\kms~for HIPASS; after degrading the LDS data to the HIPASS velocity resolution, the dashed red line is obtained.~~{\\it upper right:} \\,Comparison of the angular FWHM for the CHVCs in the HIPASS catalog  (black line) with those in the LDS catalog (dashed line). The angular resolution of the LDS is $36\\arcmin$, but $15\\arcmin$ for HIPASS; the dashed black line shows the HIPASS values after convolving with the LDS beam.~~{\\it lower left:} \\,Comparison of the total flux of CHVCs found in the HIPASS catalog (black line) with those in the LDS catalog (red line).  After compensating the HIPASS detection rates for the lower LDS sensitivity, the dashed black histogram is obtained.  The compensation was performed using the relative detection rates in the region of survey overlap (see Fig.~\\ref{fig:completeness}).~~{\\it lower right:} \\,Comparison of the total fluxes for the semi--isolated objects (:HVCs and ?HVCs) entered in the HIPASS catalog (black line) with those in the LDS listing (red line).  Compensation of the HIPASS detections to the corresponding LDS sensitivity gives the dashed black line. }\\label{fig:unity} \\end{figure*}  \\begin{figure*} \\sidecaption \\caption{Indication of the degree of completeness of the CHVC catalog extracted from the LDS by de\\,Heij et al. (\\cite{deheij02}).  The solid curve shows the fraction of external  galaxies with the indicated peak \\hi brightness temperatures that were shown by Hartmann \\& Burton (\\cite{hartmann97}) to have been detected in the LDS.  Dashed lines show the expected completeness for sensitivities that are 25\\% better or worse, respectively, than that of the LDS. The histogram indicates the fraction of HIPASS sources from the catalog of Putman et al. (\\cite{putman02a}) within each temperature range that are also found in the LDS, in the declination zone   $-30\\deg<\\delta<0\\deg$ where the two surveys overlap. }\\label{fig:completeness} \\end{figure*}  \\begin{figure*} \\centering \\caption{~Spatial deployment of CHVCs over the sky.~~{\\it upper panel:} \\,Distribution of the cataloged CHVCs, with triangles representing the LDS sample of de\\,Heij et al. (\\cite{deheij02}) at $\\delta >0\\deg$ and diamonds representing the HIPASS sample of Putman et al.  (\\cite{putman02a}) at southern declinations. Filled circles correspond to the Local Group galaxies listed by Mateo (\\cite{mateo}). Red symbols indicate positive LSR velocities and black symbols negative velocities. The background grey--scale shows \\hi column depths from an integration of observed temperatures over velocities ranging from $V_{\\rm LSR} = -450\\rm\\;km\\;s^{-1}$ to $+400\\rm\\;km\\;s^{-1}$, but excluding all gas with $V_{\\rm DEV} < 70\\rm\\;km\\;s^{-1}$.~~{\\it lower panel:} \\,Smoothed relative density field of the CHVCs, accounting for the different observational parameters of the LDS and HIPASS catalogs. The cataloged CHVCs are each represented by a Gaussian with a true--angle dispersion of $20^\\circ$; the total flux of the Gaussian is set to unity for the LDS objects and to the likelihood of observing such an object in an LDS--like survey for the HIPASS sources. The grey-scale is calibrated in object number per steradian. Contours are drawn at relative densities of $-$60\\%, $-$30\\%, 0\\%  (in white) and 30\\%, 60\\%, 90\\% (in black). A significant over--density of CHVCs in the southern hemisphere remains after accounting for the different observational parameters. }\\label{fig:skydistr} \\end{figure*}  \\begin{figure*} \\sidecaption \\caption{Kinematic deployment of CHVCs identified in the LDS (triangles) and in the HIPASS (diamonds) data, plotted against Galactic longitude for three different kinematic reference frames, namely the LSR (upper), the GSR (middle), and the LGSR (lower panel).  The filled circles show the kinematic deployment with longitude of the Local Group galaxies listed by Mateo (\\cite{mateo}).  The mean velocities and the dispersions in velocity of the CHVCs and Local Group galaxies are listed in Table~\\ref{table:velostat} for the three reference frames. }\\label{fig:glonv} \\end{figure*}  \\begin{figure*} \\sidecaption \\caption{Kinematic deployment of CHVCs identified in the LDS (triangles) and in the HIPASS (diamonds) data, plotted against Galactic latitude in the three different kinematic reference frames, as in Fig.~\\ref{fig:glonv}.  The filled circles show the kinematic deployment with latitude of the Local Group galaxies listed by Mateo (\\cite{mateo}). The mean velocities and the dispersions in velocity of the CHVCs and Local Group galaxies are listed in Table~\\ref{table:velostat} for the three reference frames. }\\label{fig:glatv} \\end{figure*}  \\begin{figure*} \\centering \\caption{Smoothed distributions of velocity and velocity dispersion of the CHVC ensemble.  The panels on the left show the average velocity in the LSR (upper), GSR (middle), and LGSR (lower) reference frames, respectively.  The panels on the right show the velocity dispersions, similarly arranged.  Individual CHVCs in the ensemble were convolved with a Gaussian of true--angle dispersion of~$20^\\circ$.  White contours for the velocity and dispersion fields are at drawn at values of $0,\\;50,\\;\\ldots\\ \\rm km\\;s^{-1}$; black ones are drawn at $-50,\\;-100,\\;\\ldots\\ \\rm km\\;s^{-1}$.  These smoothed representations of the observed situation can be compared with similarly sampled and smoothed representations of simulations, as described in the text.} \\label{fig:vfield} \\end{figure*}  \\begin{figure*} \\centering \\caption{Summary of the observed spatial, kinematic, angular size, and flux properties of the CHVC ensemble.  The three panels arranged across the top of the figure show sky projections, as follows: {\\it left:} Smoothed density field of the CHVC population.  A Gaussian with a dispersion of~$20^\\circ$ (true angle) was drawn at the location of each CHVC; the volume of the Gaussian is unity for both LDS and HIPASS sources -- thus in this case the observations are shown directly, i.e. HIPASS sources are \\emph {not} weighted by the likelihood with which they would be observed in a LDS--like survey.  {\\it middle:} Smoothed velocity field of the population in the Galactic Standard of Rest frame.  {\\it right:} Smoothed velocity dispersion field.  The grey--scale bar for the left--hand panel is labeled in units of CHVC per steradian; the other two bars are labeled in units of \\kms.  Contours are drawn at relative densities of $-$60\\%, $-$30\\%, 0\\% (in white) and 30\\%, 60\\%, 90\\% (in black). White contours for the velocity and dispersion fields are at drawn at values of $0,\\;50,\\;\\ldots\\ \\rm km\\;s^{-1}$; black ones are drawn at $-50,\\;-100,\\;\\ldots\\ \\rm km\\;s^{-1}$.  The two panels in the middle row of the figure show the kinematic distribution of the observed CHVC ensemble, representing $V_{\\rm GSR}$ plotted against $l$ and $b$, as indicated. Delta functions at the observed coordinates were convolved with a Gaussian with an angular dispersion of~$20^\\circ$ and velocity dispersion of~$20\\rm\\;km\\;s^{-1}$.  The two lower panels show, respectively, the observed peak \\hi column density distribution of the CHVC population and the observed angular size distribution. }\\label{fig:dataoverview} \\end{figure*}  \\begin{figure*} \\sidecaption \\caption{~Variation of heliocentric velocity versus the cosine of the angular distance between the solar apex and the direction of the object; CHVCs from the de\\,Heij et al. (\\cite{deheij02}) LDS compilation at $\\delta > 0\\deg$ are plotted as triangles; those from the Putman et al. (\\cite{putman02a}) HIPASS compilation, as diamonds. The CHVCs with $b~<~-65^\\circ$ are plotted in red. Local Group galaxies, from the review of Mateo~(\\cite{mateo}), are indicated by filled circles. The solid line represents the solar motion of $V_\\odot = 316\\rm\\;km\\;s^{-1}$ towards $l=93^\\circ,\\;b=-4^\\circ$ as determined by Karachentsev \\& Makarov (\\cite{karachentsev96}). Dashed lines give the  $1\\sigma(V)$~envelope ($\\pm60\\rm\\;km\\;s^{-1}$, following Sandage \\cite{sandage86}) encompassing most galaxies firmly established as members of the Local Group. } \\label{fig:vhel} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusion}  The effects of both obscuration by the gaseous disk of the Galaxy and the limited sensitivity of currently available \\hi  surveys have important consequences for the observed properties of the HVC phenomenon. We have identified those consequences in this paper. Obscuration leads to apparent localized enhancements of object density, as well as to systematic kinematic trends that need not be inherent to the population of CHVCs. A varying resolution and sensitivity over the sky substantially complicates the interpretation of the observed distributions. Taking account of both these effects in a realistic manner is crucial to assessing the viability of models for the origin and deployment of the anomalous--velocity H\\,{\\sc i}.  Our discussion leads to specific predictions for the numbers and kinematics of faint CHVCs which can be tested in future \\hi surveys.  \\subsection{Galactic Halo models} As shown in \\S\\,\\ref{sec:simplemodel}, a straightforward empirical model in which CHVCs are dispersed throughout an extended halo centered on the Galaxy does not provide the means to discriminate between distances typical of the Galactic Halo and those of the Local Group. Comparable fit quality is realized for distance dispersions ranging from about 30 to 300~kpc. In addition to requiring a relatively large number of free parameters, such empirical models beg a number of serious physical questions. In the first instance: how is it that \\hi clouds can survive at all in a low--pressure, high--radiation--density environment without the pressure support given by a dark halo? Presumably such ``naked'' Galactic Halo \\hi clouds would either be very short--lived or require continuous replenishment, since the timescales for reaching thermal and pressure equilibrium are only about 10$^7$~years (Wolfire et al.  \\cite{wolf95}). Realistic assessment of such a scenario must await more detailed simulations that track the long--term fate of gas, for example after tidal stripping from the LMC/SMC, within the Galactic Halo.  Only by including more physics will it be possible to reduce the number of free parameters and determine meaningful constraints on this type of scenario. This class of model also suffers from a number of shortcomings in describing the observed distributions, namely that the object density enhancement coupled with high negative velocities seen in the Local Group barycenter direction are not reproduced.  The Galactic Halo simulations returned formally acceptable values of characteristic distance as low as some 30 kpc.  There is, however, a growing body of independent evidence based on high--resolution imaging of a limited number of individual CHVCs that such nearby distances do not apply. Braun \\& Burton (\\cite{braun00}) discussed evidence from Westerbork synthesis observations of rotating cores in ?HVC\\,204.2\\,$+$\\,29.8\\,$+$\\,075 (using the de\\,Heij et al. \\cite{deheij02} notation for a semi-isolated source) whose internal kinematics could be well modeled by rotation curves in flattened disk systems within cold dark matter halos as parameterized by Navarro et al. (\\cite{navarro97}), if at a distance of at least several hundred kpc. Similar distances were indicated for ?HVC\\,115.4\\,$+$\\,13.4\\,$-$\\,260 on the basis of dynamical stability and crossing--time arguments regarding the several cores observed with different systemic velocities, but embedded in a common diffuse envelope.  The WSRT data for CHVC\\,125.3\\,$+$\\,41.3\\,$-$\\,205 likewise supported distances of several hundred kpc, based on a volume--density constraint stemming from the observed upper limit to the kinetic temperature of 85\\,K.  Burton et al. (\\cite{burt01}) found evidence in Arecibo imaging of ten CHVCs for exponential edge profiles of the individual objects: the outer envelopes of the CHVCs are not tidally truncated and thus are likely to lie at substantial distances from the Milky Way.  For plausible values of the thermal pressure at the core/halo interface, these edge profiles support distance estimates which range between 150 and 850 kpc.  \\subsection{Local Group models} The Local Group deployment models of \\S\\,\\ref{sec:model} offer a more self--consistent and physically motivated scenario for the CHVC population.  Dark--matter halos provide the gravitational confinement needed to produce a two--phase atomic medium with cool \\hi condensations within warm \\hi envelopes, and provide in addition the necessary protection against ram--pressure and tidal stripping to allow long--term survival. The kinematics of the population follow directly from an assumed passive evolution within the Local Group potential. While three free parameters (the distance scalelength, the mass function slope, and the upper mass cut--off) were then tuned to explore consistency with the observations, only the distance was effectively a ``free'' parameter. The mass function slopes of the best fits have values of~$-1.7$ to $-1.8$, in rough agreement with the value of $-1.6$ favored by Chiu et al. (\\cite{chiu}) for the distribution of the baryonic masses in their cosmological simulations. The somewhat steeper slopes and therefore larger baryonic fractions favored by our model fits might be accomodated by recondensation onto the dark--matter halos at later times.  The \\hi upper mass cut--off introduced in the Local Group models can also be externally constrained.  In addition to satisfying the observational demand that no \\hi column densities exceeding a few times 10$^{20}$\\,cm$^{-2}$ are seen in the CHVC population (consistent with the absence of current internal star formation), there is the observed lower limit of about 3$\\times10^7$\\,M$_\\odot$ for the \\hi mass seen in a large sample of late--type dwarf galaxies (Swaters \\cite{swat99}). The upper mass cut--off favored by the simulations, of about $10^7$\\,M$_\\odot$, is essentially unavoidable given these two constraints.  The spatial Gaussian dispersion which is favored by these simulations is quite tightly constrained to lie between about 150 and 200 kpc. The implication for the distribution of object distances is illustrated in Fig.~\\ref{fig:mod01dist} in the form of a histogram of the detected objects from model \\#9. The distribution has a broad peak extending from about 200 to 450~kpc with a few outliers extending out to 1~Mpc due primarily to the M31 sub--population.  The filled circles in the figure are the distance estimates for individual CHVCs found by Braun \\& Burton (\\cite{braun00}) and Burton et al. (\\cite{burt01}). Although very few in number, these estimates appear consistent with the model distribution, also peaking in number near 250~kpc.  We have made the simplifying assumption that the baryonic matter in our model clouds is exclusively in the form of \\hi, rather than being partially ionized. It is reassuring that the best-fitting models have peak column densities which are sufficiently high that the objects should be self-shielding to the extragalactic ionizing radiation field for M$_{\\rm HI}~>~10^{5.5}$\\,M$_\\odot$ as noted above. Since the neutral component requires a power--law slope of about $-1.7$ to fit the data, it seems likely that the total baryonic mass distribution might follow an even steeper distribution, since the mass fraction of ionized gas will increase toward lower masses.  \\subsection{The Local Group mass function} An interesting question to consider is whether the extrapolated mass distributions of our Local Group CHVC models can also account for the number of galaxies currently seen. In Fig.~\\ref{fig:lgall} we plot the mass distribution of objects in one of the best--fitting Local Group models, model \\#9 of Table~\\ref{table:bestfit}. The thin--line histogram gives the mass distribution of the model population after accounting for the effects of ram--pressure and tidal stripping. The thick--line histogram gives the observed CHVC distribution that results from applying the effects of Galactic obscuration and sensitivity limitations appropriate to the LDS and HIPASS properties in the northern and southern hemispheres, respectively. The hatched histogram gives the inferred total baryonic (\\hi plus stellar) mass distribution of the Local Group galaxies tabulated by Mateo (\\cite{mateo}), assuming a stellar mass--to--light ratio of $M/L_B = 3$\\,M$_\\odot$/L$_\\odot$. M31 and the Galaxy, with baryonic masses of some 10$^{11}$\\,M$_\\odot$, are not included in the plot.  The diagonal line in the figure has the slope of the model \\hi mass function of $\\beta=-1.7$. The figure demonstrates that the low--mass populations of these models are roughly in keeping with what is expected from the number of massive galaxies together with a constant mass function slope of about $\\beta=-1.7$. At intermediate masses, 10$^7$--10$^{8.5}$\\,M$_\\odot$, there is a small deficit of cataloged Local Group objects relative to this extrapolated distribution, while at higher masses there is a small excess. Conceivably this may be the result of galaxy evolution by mergers.  It is important to note that the distribution of objects shown in Fig.~\\ref{fig:lgall} is only the current relic of a much more extensive parent population. As shown in Table~\\ref{table:modstat}, about 75\\% of the CHVC population in these models is predicted to have been disrupted by ram pressure or tidal stripping over a Hubble time, contributing about $3\\times10^9$~M$\\odot$ of baryons to the Local Group environment and the major galaxies.  \\subsection{The M\\,31 population of CHVCs} One of the most suggestive attributes of the CHVC population in favor of a Local Group deployment is the modest concentration of objects which are currently detected in the general direction of M\\,31, i.e. in the direction of the Local Group barycenter.  These objects have extreme negative velocities in the GSR reference frame. While this is a natural consequence of the Local Group models it does not follow from the empirical Galactic halo models, nor is it a consequence of obscuration by Galactic H\\,{\\sc i}. Putman \\& Moore (\\cite{putman02c}) have made some comparisons between numerical simulations of dark matter mini--halos in the Local Group with the $(l,V_{\\rm LGSR})$ distributions of HVCs and CHVCs, and were led to reject the possibility of CHVC deployment throughout the Local Group. Our discussion here has shown that such comparisons require taking explicit account of detection thresholds in the available survey observations, as well as of the vagaries of obscuration caused by the \\hi Zone of Avoidance. The Putman \\& Moore investigation did not take these matters into account. The modest apparent amplitude of the M\\,31 concentration relative to the Galactic population as seen with present survey sensitivities provides the best current constraints on the global distance scale of the CHVC ensemble. There follows a testable prediction, namely that with increased sensitivity a larger fraction of the M\\,31 population of CHVCs should be detected. This prediction was made explicit in Fig.~\\ref{fig:mod01hom}, where one of our model distributions was shown as it would have been detected if HIPASS sensitivity were available in the northern sky. For that particular model, some 250 additional detected objects are predicted, of which the majority are concentrated in the $60\\times60^\\circ$ region centered on M\\,31. The ongoing HIJASS survey of the sky north of $\\delta=25^\\circ$ (Kilborn \\cite{kilborn}), which is being carried out using the 76--m Lovell Telescope at Jodrell Bank to about the same velocity coverage, angular resolution, and sensitivity as the HIPASS effort, should allow this prediction to be tested.  \\subsection{The Sculptor Group lines of sight} We have omitted the part of the sky around the south Galactic pole in our fitting of Local Group models to the observations, because of the extreme velocity dispersions measured in this direction. The nearest external group of galaxies, the Sculptor Group, is located in the direction of the south Galactic pole.  If the CHVCs are distributed around the major Local Group galaxies, then plausibly the same sort of objects could be present in the Sculptor Group.  Putman et al. (\\cite{putman02a}) mention detection of clouds in the direction of the southern part of the Sculptor Group.  Because no similar clouds were detected in the northern part of this Group, they consider it unlikely that this concentration of CHVCs is associated with the Sculptor Group. We note, however, that rather than being a spherical concentration of galaxies, the Sculptor Group has an extended filamentary morphology, which ranges in distance from~$1.7\\rm\\;Mpc$ in the south to~$4.4\\rm\\;Mpc$ in the north.  Putman et al. assumed that the HIPASS sensitivity would allow detection of \\hi masses of~$7\\times10^6\\rm\\;M_\\odot$ throughout the Sculptor Group.  But in Fig.~\\ref{fig:obsdistance} we show the actual distance out to which HIPASS can detect \\hi masses given a realistic cloud model and detection threshold: even the most massive and rare objects in our simulated distributions, with $M_{\\rm HI}$~=~$10^7\\rm\\;M_\\odot$, can only be detected out to $2.5\\rm\\;Mpc$. It is therefore only the near portion of the Sculptor filament that might be expected to show any enhancement in CHVC density with the currently available sensitivities.  \\subsection{Predicted CHVC populations in other galaxy groups} It is also interesting to consider whether the simulated Local Group model populations would be observable in external galaxy groups at even larger distances. In Fig.~\\ref{fig:mod01mass} we show one of our best--fitting Local Group models, model \\#9 of Table~\\ref{table:bestfit}, projected onto a plane as in Fig.~\\ref{fig:mod3d}. In Fig.~\\ref{fig:mod3d}, the surviving clouds were distinguished by \\hi flux; in Fig.~\\ref{fig:mod01mass}, the distinction is by \\hi mass.  We indicate with grey dots those objects that were deemed to have been disrupted by ram--pressure or tidal stripping. The red and black dots indicate the remaining objects in the population, with the red dots representing objects that exceed $M_{\\rm HI}$~=~3$\\times10^6$\\,M$_\\odot$ and the black dots those that fall below this mass limit. The choice of a limiting mass of $M_{\\rm HI}$=3$\\times10^6$\\,M$_\\odot$ over a linewidth of 35~\\kms~was made to represent what might be possible for a deep \\hi survey of an external galaxy group. In this example, some 95 objects occur which exceed this mass limit distributed over a region of some 1.5$\\times$1.0 Mpc extent. For a limiting mass of $M_{\\rm HI}$=5$\\times10^6$\\,M$_\\odot$ over 35~\\kms, the number drops to 45. It is clear that a very good mass sensitivity will be essential to detecting such potential CHVC populations in external galaxy groups. Current searches for such populations, reviewed by Braun \\& Burton (\\cite{braun01}), have generally not reached a sensitivity as good as even $M_{\\rm HI}$=$10^7\\rm\\;M_\\odot$ over 35~\\kms, so it is no surprise that such distant CHVCs have not yet been detected."
    },
    "0206/astro-ph0206130_arXiv.txt": {
        "abstract": "Subdwarf B (sdB) stars (and related sdO/sdOB stars) are believed to be helium core-burning objects with very thin hydrogen-rich envelopes. In recent years it has become increasingly clear from observational surveys that a large fraction of these objects are members of binary systems.  To better understand their formation, we here present the results of a detailed investigation of the three main binary evolution channels that can lead to the formation of sdB stars: the common envelope (CE) ejection channel, the stable Roche lobe overflow (RLOF) channel and the double helium white dwarfs (WDs) merger channel. The CE ejection channel leads to the formation of sdB stars in short-period binaries with typical orbital periods between 0.1 and 10\\,d, very thin hydrogen-rich envelopes and a mass distribution sharply peaked around $\\sim 0.46M_\\odot$. On the other hand, under the assumption that all mass transferred is soon lost, the stable RLOF channel produces sdB stars with similar masses but long orbital periods (400 -- 1500\\,d) and with rather thick hydrogen-rich envelopes.  The merger channel gives rise to single sdB stars whose hydrogen-rich envelopes are extremely thin but which have a fairly wide distribution of masses (0.4 -- 0.65\\,$M_{\\odot}$).  We obtained the conditions for the formation of sdB stars from each of these channels using detailed stellar and binary evolution calculations where we modelled the detailed evolution of sdB stars and carried out simplified binary population synthesis simulations.  The observed period distribution of sdB stars in compact binaries strongly constrains the CE ejection parameters. The best fits to the observations are obtained for very efficient CE ejection where the envelope ionization energy is included, consistent with previous results. We also present the distribution of sdB stars in the $T_{\\rm eff}$ - $\\log g$ diagram, the Hertzsprung-Russell diagram and the distribution of mass functions. ",
        "introduction": "Subdwarf B (sdB) stars were originally defined by Sargent and Searle \\shortcite{sar68} as stars with colours corresponding to those of B stars in which the Balmer lines are abnormally broad compared to those seen in population I main-sequence stars. Subdwarf O (sdO) stars and subdwarf OB (sdOB) stars are related stars of correspondingly earlier spectral type (see, e.g., Vauclair \\& Liebert 1987).  Based on an interpretation of their evolutionary state, sdB stars are also sometimes referred to as extreme horizontal branch stars. They are generally considered to be core helium burning stars with extremely thin hydrogen envelopes ($<0.02M_\\odot$), and most of them are believed to have masses around $0.5M_\\odot$ \\cite{heb86,saf94}. Indeed, a recent asteroseismological analysis by Brassard et al.\\ (2001) has confirmed a mass of 0.49$\\pm 0.02M_{\\odot}$ for the sdB star PG 0014+067.  In this paper, we collectively refer to core helium burning stars with thin hydrogen envelopes as sdB stars if they are located in the corresponding region in a ($T_{\\rm eff}$, $\\log g$) diagram, even if some of them may in reality be sdO or sdOB stars.    Subdwarf B stars form an important class of objects in several respects.  At the Galactic level, they are the dominant population in surveys of blue objects \\cite{gre86} and constitute a population of stars that are important for our understanding of the structure and evolution of the Galaxy. Pulsating sdB stars \\cite{kil99} provide a standard candle for distance determinations. On a larger cosmological scale they have been used to constrain the ages of the oldest galaxies and hence cosmological models. The latter is based on measuring the age of giant elliptical galaxies from the ultraviolet (UV) excess, or ``upturn'', with the help of evolutionary population synthesis models where low mass core-helium burning stars provide the dominant source of UV radiation \\cite{bro97,yi97,yi99}.   More importantly, sdB stars are exotic objects because of their thin hydrogen-rich envelopes. Understanding the process of their formation helps to improve our understanding of the theory of stellar and binary evolution.  There have been extensive surveys of sdB stars in the past. Magnitude-limited and colour-selected samples have been obtained from the Palomar Green (PG) survey \\cite{gre86} (magnitude limit $B\\sim 16.1$) and the Kitt Peak Downes (KPD) survey \\cite{dow86} (magnitude limit $B=15.3$).  Saffer et al.\\ (1994) measured atmospheric parameters, such as effective temperature, surface gravity and photospheric helium abundance, for 68 sdB stars.  Ferguson, Green \\& Liebert \\shortcite{fer84} found 19 sdB stars with main sequence (MS) companions from the PG survey and derived a binary frequency of about 50 percent. Allard et al.\\ \\shortcite{all94} found 31 sdB binaries from 100 candidates chosen from the PG and the KPD surveys and estimated that 54 to 66 percent of sdB stars are in binaries with MS companions after taking selection effects into account.  Thejll, Ulla \\& MacDonald \\shortcite{the95} and Ulla \\& Thejll \\shortcite{ull98} also found that more than half of their sdB star candidates showed infrared flux excesses, indicating the presence of binary companions. Aznar Cuadrado \\& Jeffery \\shortcite{azn01} obtained atmospheric parameters for 34 sdB stars from spectral energy distributions and found that 19 were binaries with MS companions, while 15 appeared to be single.  These observations showed that at least half of the sdB stars were in binaries.  A major recent development has been the identification of many sdB stars as short-period binaries \\cite{saf98,koe98,jef98,woo99,oro99,mor99,max00a,max00b,max01,heb02}. In particular, Maxted et al.\\ \\shortcite{max01} concluded that more than two thirds of their candidates were binaries with short orbital periods from hours to days and that most of the known companions were white dwarfs (WDs).  A variety of formation channels for sdB stars have been proposed in the past but mainly for single sdB stars because of the absence of identified sdB star binaries at the time.  In the merger channel, two helium white dwarfs in a close binary are driven together by the orbital angular momentum loss due to gravitational wave radiation. When the white dwarfs merge and the merged object ignites helium, this produces a single sdB star \\cite{web84,ibe86,han98}.  Alternatively, stellar wind mass loss near the tip of the first giant branch (FGB) may strip off a giant's envelope and leave an almost bare helium core.  If helium is ignited in the core, the star will appear as a single sdB star \\cite{dcr96}. Sweigart \\shortcite{swe97} has studied the evolution of globular-cluster stars and suggested that helium mixing driven by internal rotation substantially increases the helium abundance in the envelope; this may lead to enhanced mass loss along the FGB and the formation of a sdB star.   On the other hand, Mengel, Norris \\& Gross \\shortcite{men76} carried out conservative binary evolution calculations for a binary system with initial masses of $0.80M_\\odot$ and $0.78M_\\odot$ and a composition $X=0.73$, $Z=0.001$, and showed that there exists a range of initial separations for which stable mass transfer can produce an sdB star of $\\sim 0.5M_\\odot$ in a wide binary.  From a binary evolution point of view, these formation channels are not complete.  When a star fills its Roche lobe near the tip of the FGB, mass transfer begins and may be dynamically unstable.  This leads to the formation of a common envelope (CE) \\cite{pac76}, where the CE engulfs the helium core and the secondary.  Due to friction between the envelope and the immersed binary, the orbit shrinks, depositing a large amount of orbital energy in the envelope. If this energy is enough to eject the envelope and if helium is subsequently ignited in the core, a sdB star in a short-period binary is formed with a mass near $0.5M_\\odot$. These are exactly the types of objects identified in large numbers by Maxted et al.\\ \\shortcite{max01}.  If mass transfer near the tip of the FGB is dynamically stable, the envelope of the primary is lost as a result of stable RLOF, and the remnant core will be in a binary system with a long orbital period. It becomes a sdB star when helium in the primary's remnant is ignited. An additional channel for the formation of sdB stars in wide binaries, which has not received much attention in the past, involves binaries that experience stable RLOF when passing through the Hertzsprung gap (so-called early case B mass transfer) (Han, Tout \\& Eggleton 2000; Han et al. 2002; in preparation [henceforth, Paper II]).  All of the sdB binaries produced through stable RLOF channels are consistent with the observations by Green, Liebert \\& Saffer \\shortcite{gre00} who showed that some sdB stars appear to be members of long-period binaries.   The main purpose of this study is to re-examine the various scenarios for the formation of sdB stars in some detail. In this first paper, we concentrate on the individual evolutionary channels. Using detailed stellar and binary calculations, we model the physics and appearance of sdB stars and then test individual evolutionary channels using binary population synthesis (BPS). We demonstrate that all of the main evolutionary channels proposed previously can lead to the formation of sdB stars. As a by-product we constrain the CE ejection efficiency from the observed period distribution of compact sdB binaries to arrive at a physically motivated and experimentally calibrated prescription for the CE phase.   The outline of this paper is as follows.  In section 2, we describe the stellar evolution code and the binary population synthesis code adopted in this study. In section 3, we present the conditions for the formation of sdB stars from the CE ejection channel, their evolutionary tracks and simplified BPS models to constrain the CE ejection efficiency.  In section 4, we derive the conditions for helium ignition in objects that result from the merger of two He white dwarfs and use Monte Carlo simulations to determine their mass distribution.  In section 5, we investigate the criterion for stable RLOF and the formation of sdB stars in wide binaries.  In the follow-up paper (Paper II) we will apply these results to a comprehensive binary population synthesis study and will estimate the relative importance of these individual channels. ",
        "conclusions": "In this paper, we have demonstrated that the three binary evolution channels that have been proposed for the formation of sdB (and related sdO/sdOB) stars may all contribute to the observed population.  In the CE ejection channel, which may account for more than 2/3 of all sdB stars, dynamically unstable mass transfer near the tip of the FGB results in the formation of a CE and spiral-in phase, leaving a short-period binary after the envelope has been ejected. The system becomes a sdB binary if helium is ignited. Using detailed stellar evolution calculations, we have determined how close to the tip of the FGB the progenitor has to be at the onset of RLOF. Using simplified binary population synthesis calculations, we have been able to show that the CE ejection process has to be very efficient and that the ionization energy in the envelope has to be included in the ejection criterion in order to be able to explain the observed orbital period distribution.  In the stable RLOF channel, the progenitor systems experiences stable mass transfer in which the giant is stripped off its envelope as a result of the mass transfer. If this occurs near the tip of the FGB, the remnant helium core will still ignite helium in the core and become a sdB star in a binary with a long orbital period and a fairly thick hydrogen-rich envelope as compared to the other channels. Using detailed binary evolution calculations, we demonstrated that this channel usually requires a fairly massive white dwarf companion or enhanced stellar wind mass loss before the onset of RLOF (e.g. tidally enhanced winds).  Double He WDs may coalesce due to gravitational wave radiation. When helium is ignited in the merger, a single sdB star is formed, and its hydrogen envelope is likely to be very thin. We have determined the conditions for which the merged system will be able to ignite helium.  In a follow-up paper we will implement these results in full binary population synthesis calculations to assess their relative importance and to allow direct comparison with observed subdwarf populations."
    },
    "0206/astro-ph0206289_arXiv.txt": {
        "abstract": "We have detected a population of predominantly blue $(B-V \\le 1.1)$ stars in the direction $l = 167\\arcdeg, b = -35\\arcdeg$ (Kapteyn Selected Area 71) that cannot be accounted for by standard starcount models. Down to $V \\sim 20$, the colors and magnitudes of these stars are similar to those of the southern overdensity detected by the Sloan Digital Sky Survey at $l = 167\\arcdeg, b = -54\\arcdeg$, and identified as stripped material from the Sagittarius dwarf spheroidal galaxy. We present absolute proper motions of the stars in SA 71, and we find that the excess blue stars represent a distinct, kinematically cooler component than the Galactic field, and in reasonable agreement with predictions of Sgr disruption models. The density of the excess SA 71 stars at $V \\sim 18.8$ and $B-V \\le 1.1$ is within a factor of two of the density of the SDSS-south Sgr stripped material, and of that predicted by the Helmi \\& White disruption model. Three additional anticenter fields (SA 29, 45 and 118) show very good agreement with standard starcount models. ",
        "introduction": "The Sagittarius dwarf galaxy (Sgr) (Ibata, Gilmore \\& Irwin 1994) has drawn considerable interest as a ``living'' example of the accretion process supposedly responsible for the creation of the outer Galactic halo. Most early Sgr surveys focused within $10$-$15\\arcdeg$ of its core (e.g., Ibata {\\it et al.} 1997, hereafter I97, Mateo {\\it et al.} 1998). More recently, Sgr-associated material has been identified farther from its main body (e.g., Majewski {\\it et al.} 1999, hereafter M99), and it is becoming apparent that Sgr tidal debris may envelop the whole Galaxy roughly along the satellite's orbit. Results from the Sloan Digital Sky Survey (SDSS) in two equatorial strips (Yanny {\\it et al.} 2000, hereafter Y00; Ivezic {\\it et al.} 2000, hereafter I00) find RR Lyrae stars, blue horizontal branch (BHB) stars, and/or blue stragglers (BS) in excess of the Galactic population that are most plausibly explained as Sgr material (see Ibata {\\it et al.} 2001a). Newberg {\\it et al.} (2002, hereafter N02) confirm the previous SDSS Sgr detections with color-magnitude diagrams (CMDs) showing main sequence turnoff colors (MSTO) consistent with those of Sgr. Similarly, Mart\\'{i}nez-Delgado {\\it et al.} (2001a,b; collectively referred to as M01) report Sgr detections associated with MSTO stars in several regions along the Sgr orbital path. Vivas {\\it et al.} (2001) confirm the I00 detection of RR Lyrae stars in a $\\sim 50$ kpc distant clump at positive Galactic latitudes, while Dohm-Palmer {\\it et al.} (2001) and Kundu {\\it et al.} (2002; K02 hereafter) report giant star concentrations in the inner halo having radial velocities (RVs) and distances in good agreement with wrapped tidal arms predicted from Sgr disruption models. Ibata {\\it et al.} (2001b) also show that the carbon-star distribution from the Totten \\& Irwin (1998, hereafter TI98) all-sky survey is highly correlated with the orbital path of Sgr.  Here we present absolute proper motions and $B-V$ color  distributions in Kapteyn Selected Area (SA) 71, a Galactic anticenter field located close to the orbital path of Sgr ($ l = 167 \\arcdeg, b = -35 \\arcdeg$). We find a population of predominantly blue stars ($ B-V \\le 1.1$) in excess of that predicted by standard Galactic starcount models. The proper motions of the excess blue stars show a distinct, kinematically cooler population than the Galactic field. Assuming that the excess blue stars are horizontal branch (HB) stars, we estimate their distance to be $\\sim 30$ kpc, a value consistent with Sgr disruption model predictions (e.g., Johnston 1998, hereafter J98). We also present $B-V$ color distributions for three other anticenter fields (SA 29, 45 and 118). These allow us to refine some of the parameters of the starcount model. These three fields are consistent with each other and agree with the global model predictions. This work is a continuation of the photometric and proper motion study of SA 57 by Majewski (1992) as well as the Mount Wilson Halo Mapping Project (Sandage 1997). ",
        "conclusions": ""
    },
    "0206/astro-ph0206240_arXiv.txt": {
        "abstract": "{ We consider the structure of steady--state radiative shock waves propagating in the partially ionized hydrogen gas with density $\\rho_1 = 10^{-10}~\\gcc$ and temperature $3000\\K\\le T_1\\le 8000\\K$. The radiative shock wave models with electron thermal conduction in the vicinity of the viscous jump are compared with pure radiative models. The threshold shock wave velocity above of which effects of electron thermal conduction become perceptible is found to be of $U_1^*\\approx 70~\\kms$ and corresponds to the upstream Mach numbers from $M_1\\approx 6$ at $T_1=8000\\K$ to $M_1\\approx 11$ at $T_1=3000\\K$. In shocks with efficient electron heat conduction more than a half of hydrogen atoms is ionized in the radiative precursor, whereas behind the viscous jump the hydrogen gas undergoes the full ionization. The existence of the electron heat conduction precursor leads to the enhancement of the Lyman continuum flux trapped in the surroundings of the discontinuous jump. As a result, the partially ionized hydrogen gas of the radiative precursor undergoes an additional ionization ($\\delta\\xH\\lesssim 5\\%$), whereas the total radiave flux emerging from the shock wave increases by $10\\%\\le\\delta(\\FR)\\le 25\\%$ for $70~\\kms\\le U_1 \\le 85~\\kms$. ",
        "introduction": "In our previous papers \\citep[][]{1998A&A...333..687F,2000A&A...354..349F,2001A&A...368..901F}, hereinafter referred to as Papers~I -- III, we presented the studies of steady--state radiative shock waves propagating in the partially ionized hydrogen gas with properties that are typical for atmospheres of pulsating late--type stars. The models were considered in terms of the self--consistent solution of the equations of fluid dynamics, radiation transfer and rate equations for the hydrogen atom and provide the reliable estimates of radiative energy losses of the shock wave. In these studies we adopted that at the viscous jump which is treated as an infinitesimally thin discontinuity the electron gas compresses adiabatically. Such an assumption is justified by the very low rate of energy exchange between protons and electrons within the viscous jump and thereby allows us to determine the postshock electron temperature from the simple adiabatic relation. However, it is known that the characteristic length scale of electron thermal conduction might be comparable with the thickness of the postshock relaxation zone (the length of the temperature equilibration zone) and, therefore, might affect perceptibly the spatial distribution of hydrodynamic variables at least in the vicinity of the viscous jump \\citep[see, for discussion,][]{1967pswh.book.....Z,1984frh..book.....M}.  Effects of electron thermal conduction in two--temperature shock waves propagating in partially ionized helium and argon with temperatures and densities close to those of stellar atmospheres were investigated using both hydrodynamic and kinetic approaches by \\cite{Grewal:Talbot:1963,Jaffrin:1965,Lu:Huang:1974}. Solution of the more general problem involving the radiation transfer was considered by \\cite{Vinolo:Clarke:1973}. The conspicuous result of these calculations is that the authors demonstrated the existence of the zone of the elevated electron temperature ahead the viscous jump appearing due to the high thermal conductivity of the electron gas. For the hydrogen gas effects of electron thermal conduction were considered only in the limit of full ionization at Mach numbers as high as $M_1\\approx 8$ \\citep{1993PhFlB...5.3182V,1995PhPl....2.1412V}. This study, unfortunately, was confined to the problem of inertial confinement fusion for the extremely high temperature and number density of electrons ($T = 10^8\\K$, $\\nel = 10^{22}~\\cm^{-3}$).  By now, the scarce studies of astrophysical radiative shock waves with electron thermal conduction in two--temperature gases were confined to the interstellar medium \\citep{1989ApJ...336..979B,1990ApJ...348..169B} and accreting white dwarfs \\citep{1987ApJ...313..298I}. In these works the heat conduction was found to substantially affect the structure of the postshock relaxation zone for shocks with velocities ranged from $70~\\kms$ to $170~\\kms$. In atmospheres of late--type pulsating stars the gas density is many orders of magnitude higher in comparison with that of the interstellar medium, whereas the velocities of shocks does not exceed $100~\\kms$, so that the role of electron heat conduction in radiative losses remains highly uncertain. In this paper we compute the models of steady--state radiative shock waves propagating through the partially ionized hydrogen gas with properties typical for atmospheres of pulsating late--type stars and compare the pure radiative shock wave models with those in which electron the thermal conduction is taken into account.  As in our previous papers we assume that the ambient unperturbed medium is homogeneous and effects of magnetic fields are negligible. Indeed, the magnetic field becomes important when the electron mean free path $\\mfp$ is larger than the gyromagnetic radius $\\gyr$. For the gas density considered in our study ($\\rho = 10^{-10}~\\gcc$) the condition $\\mfp\\gg\\gyr$ fulfills for $B\\gg 3$~Gauss. Furthermore, the magnetic pressure $B^2/8\\pi$ becomes comparable with the gas pressure for $B > 25$~Gauss. On the other hand, the strength of the magnetic field in late--type giants is $B\\ll 1$~Gauss \\citep{1982A&A...105..133B}. Thus, effects of magnetic fields can be ignored without any loss of accuracy. We expect that results of our calculations can be applied to shock phenomena observed in atmospheres of radially pulsating giants and supergiants such as W~Vir, RV~Tau and Mira type variables.  In order to specify the model we use three parameters determining the structure of the steady--state radiative shock wave. These are the density $\\rho_1$ and the temperature $T_1$ of the ambient hydrogen gas and the upstream gas flow velocity $U_1$. ",
        "conclusions": "In this paper we have shown that in radiative shock waves with upstream velocities of $U_1\\gtrsim 70~\\kms$ the physical properties of the gas surrounding the discontinuous jump are sufficient for appearence of efficient electron heat conduction, the conductive flux being comparable with the Lyman continuum flux in the vicinity of the discontinuous jump. The existence of the narrow conductive precursor affects all the region of the Lyman continuum radiation trapped around the discontinuous jump between zones of preshock ionization and postshock recombination of the hydrogen gas. As a consequence, the hydrogen ionization degree increases in the radiative precursor by as much as 5\\% provided that the preshock hydrogen gas is partially ionized. The effect of the saturation of the electron conductive flux can be important only ahead the discontinuity but for models considered in the present study we obtained $\\vert\\Fe\\vert\\lesssim\\Fsat$. The main conclusion of our study is that in shocks with velocities exceeding the threshold value $U_1^*\\approx 70~\\kms$ the electron thermal conduction significantly raises the radiative losses and, therefore, can diminish the efficiency of the shock--driven mass loss. More studies of radiative shock waves with electron heat conduction for wider ranges of $\\rho_1$, $T_1$ and $U_1$ are needed."
    },
    "0206/astro-ph0206076_arXiv.txt": {
        "abstract": "In this paper we examine various issues closely related to the ongoing discussion on the nature of the Blandford-Znajek mechanism of extraction of rotational energy of black holes. In particular, we show that switch-on and switch-off shocks are allowed by the shock equations of relativistic MHD and have similar properties to their Newtonian counterparts. Just like in Newtonian MHD they are limits of fast and slow shock solutions and as such they may be classified as weakly evolutionary shocks. The analysis of Punsly's MHD waveguide problem shows that its solution cannot have the form of a traveling step wave and that both fast and Alfv\\'en waves are essential for generating the flow in the guide. Causality considerations are used to argue that the Blandford-Znajek perturbative solution is in conflict with the membrane paradigm. An alternative interpretation is presented according to which the role of an effective unipolar inductor in the Blandford-Znajek mechanism is played by the ergospheric region of a rotating black hole. Various implications of this are discussed. ",
        "introduction": "In coordinate systems singular at the black hole horizon, like the popular Boyer-Lindquist coordinates, the horizon is inevitably turns into a rather peculiar boundary of spacial domain. According to the widely accepted ``Membrane paradigm'' the horizon, or rather somewhat less stringently defined ``stretched horizon'' which is placed somewhere just above the real horizon, may be identified with a rotating conducting sphere (e.g. Blandford 1979, Thorne et al. 1986). This makes magnetized black holes look analogous to magnetized neutron stars. For many years, beginning with \\cite{PC}, Brian Punsly have been criticizing this view and also the perturbative steady-state electromagnetic wind solution for a force-free magnetosphere of a rotating black hole due to Blandford and Znajek \\shortcite{BZ77}, the BZ solution, together with similar MHD models (e.g. Phinney 1982,1983) on the basis of causality arguments. Indeed, in the case of pulsars there is only an outgoing wind which passes first through the Alfv\\'en critical surface and then through the fast critical surface and, thus, the neutron star can communicate with the wind by means of both fast and Alfv\\'en waves (Since the gas pressure is dynamically insignificant in the tenuous magnetospheres of neutron stars and black holes, the slow waves seem to be irrelevant.)  Black holes, however, must also develop an ingoing wind which passes through its own pair of critical surfaces before reaching the horizon \\cite{Jap90}. For the typical parameters of astrophysical black holes the inner fast surface is likely to be extremely close to the black hole horizon and one may argue that the stretched horizon can communicate with the outgoing wind by means of fast waves. On the other hand, the inner critical Alfv\\'en surface may be rather distant from the black hole horizon. Thus, even the stretched horizon cannot communicate with the outgoing wind by means of Alfv\\'en waves which makes black holes rather different from neutron stars.  Blandford (1979) proposed that the outgoing wind of black holes is established by means of fast waves alone (This seems to be the only way to reconcile the membrane paradigm with the BZ solution.)  However, Punsly (1996,2001) argued that fast waves are completely irrelevant and Alfv\\'en waves are solely responsible for creating the global system of poloidal electric currents of such winds and adjusting the angular velocity of magnetic field lines  and  suggested that the steady-state BZ solution is unstable and, hence, nonphysical. In particular, Punsly criticized Znajek's horizon boundary condition \\cite{Z77} used to determine the wind constants in the BZ electrodynamic solution and in the MHD analysis of Phinney (1982,1983) as a condition imposed in a region causally disconnected from the outgoing wind.  However, Znajek's condition is not really a boundary condition as it simply prohibits infinitely strong electromagnetic field as measured by a physical observer, e.g. a free falling observer (e.g. Phinney 1983). Recently, Beskin and Kuznetsova \\shortcite{BK00} stressed this point once more and argued that, though the interpretation of the stretched horizon as a unipolar inductor is misleading, there is nothing wrong with causality of the BZ-like MHD models.   This conclusion is strongly supported by the results of recent time-dependent electrodynamic simulations \\cite{Kom01} which indicate asymptotic stability of the BZ solution (Znajek's boundary condition was not imposed in these simulations.)  In this paper we continue the discussion of the causality paradox a bit further and attempt to clarify the nature of BZ mechanism. At first sight, the study presented here may appear rather unfocused but a closer look reveals strong connections between its sections. In Sec.2 we study limit shocks of relativistic MHD. In Sec.3 we use these results to analyse the Punsly's waveguide problem that provides important insights into the problem of relativistic MHD and electrodynamic winds. In Sec.4 we propose a modification of the membrane paradigm that does not conflict with causality. ",
        "conclusions": "\\begin{enumerate} \\item Switch-on and switch-off shocks are allowed by the shock equations of relativistic MHD and have similar properties to their Newtonian counterparts. Just like in Newtonian MHD they are limits of fast and slow shock solutions and as such they may be classified as weakly evolutionary shocks.  \\item Contrary to what is claimed in \\cite{P01}, the solution to Punsly's MHD waveguide problem cannot have the form of a step-like traveling wave and the guide flow cannot be established by means of Alfv\\'en waves alone.  Even in the limit of degenerate electrodynamics where a step-wave solution exists it involves a mixture both fast and Alfv\\'en waves. This suggests that both waves are important in the problems of magnetically driven MHD and electrodynamic winds.  \\item Blandford-Znajek solution contradicts to the membrane paradigm as the stretched horizon cannot play the role of a unipolar inductor. Causality arguments suggest that, just like in the case of the Penrose mechanism, the driving ``force'' of the Blandford-Znajek mechanism is the ergospheric region of space-time. \\end{enumerate}"
    },
    "0206/astro-ph0206183_arXiv.txt": {
        "abstract": "We report the discovery and monitoring of radio emission from the Type Ic \\sn{} ranging in frequency from 1.43 to 22.5 GHz, and in time from 4 to 50 days after the SN explosion.  As in most other radio SNe, the radio spectrum of \\sn{} shows evidence for absorption at low frequencies, usually attributed to synchrotron self-absorption or free-free absorption. While it is difficult to discriminate between these two processes based on a goodness-of-fit, the {\\it unabsorbed} emission in the free-free model requires an unreasonably large ejecta energy.  Therefore, on physical grounds we favor the synchrotron self-absorption (SSA) model.  In the SSA framework, at about day 2, the shock speed is $\\approx 0.3c$, the energy in relativistic electrons and magnetic fields is $\\approx 1.5\\times 10^{45}$ erg and the inferred progenitor mass loss rate is $\\approx 5\\times 10^{-7}$ M$_\\odot$/yr (assuming a $10^3$ km sec$^{-1}$ wind).  These properties are consistent with a model in which the outer, high velocity supernova ejecta interact with the progenitor wind.  The amount of relativistic ejecta in this model is small, so that the presence of broad lines in the spectrum of a Type Ib/c supernova, as observed in SN\\,2002ap, is not a reliable indicator of a connection to relativistic ejecta and hence $\\gamma$-ray emission. ",
        "introduction": "\\label{sec:intro}  Type Ib/c supernovae (SNe) enjoyed a broadening in interest over the last few years since their compact progenitors (Helium or Carbon stars) are ideal for detecting the signatures of a central engine.  Such an engine is expected in the collapsar model (Woosley 1993; MacFadyen, Woosley \\& Heger 2001)\\nocite{woo93,mwh01}, the currently popular model for long-duration $\\gamma$-ray bursts (GRBs).  In this model, the engine (a rotating and accreting black hole) provides the dominant source of explosive power.  The absence of an extensive Hydrogen envelope in the progenitor star may allow the jets from the central engine to propagate to the surface and subsequently power bursts of $\\gamma$-rays.  Separately, the Type Ic SN\\,1998bw (Galama et al. 1998)\\nocite{gvv+98} found in the localization region of GRB\\,980425 (Pian et al. 2000)\\nocite{paa+00} ignited interest in ``hypernovae''\\footnotemark\\footnotetext{There is no accepted definition for a hypernova. Here we use the term to mean a supernova with an explosion energy significantly larger than $10^{51}$ erg.}. Regardless of the controversy over the association of SN\\,1998bw with GRB\\,980425, or equivalently, the controversy over the relation of the extremely underluminous GRB\\,980425 to the cosmological GRBs, one fact is not in dispute: {\\it SN\\,1998bw is a most interesting SN.} First, the SN exhibited the broadest absorption lines to date, about 60,000 km sec$^{-1}$ (Iwamoto et al. 1998; Woosley, Eastman \\&\\ Schmidt 1999)\\nocite{imn+98,wes99}.  Second, modeling of the optical spectra and lightcurves suggested a large energy release, $E_{51}\\sim 30$, where $E_{51}$ is the SN energy release in units of $10^{51}$ erg. Finally, and most relevant to the issue of GRB connection, the SN was the brightest radio SN at early times; robust equipartition arguments led to an inferred energy of $E_{\\Gamma}\\gtrsim 10^{49}$ erg in ejecta with relativistic velocities, $\\Gamma\\sim$ few (Kulkarni et al. 1998, hereafter K98).  Until SN\\,1998bw, no other SN showed hints of such an abundance of relativistic ejecta. Tan, Matzner \\&\\ McKee (2001)\\nocite{tmm01} explain the relativistic ejecta as resulting from an energetic shock as it speeds up the steep density gradient of the progenitor. The $\\gamma$-ray and radio emission would then arise in the forward shock.   From the perspective of a GRB--SN connection, what matters most is the presence of relativistic ejecta. $\\gamma$-ray emission traces ultra-relativistic ejecta, but as was dramatically demonstrated by SN\\,1998bw, the radio serves as an equally good proxy for relativistic ejecta with the added advantage that the emission is not beamed. Given this, we began a systematic program of investigating at radio wavelengths all Ib/c SNe with features similar to SN\\,1998bw: a hypernova or broad optical lines.  Y. Hirose discovered \\sn{} in M74 (distance, $d\\sim 7.3\\,$Mpc; Smartt et al.  2002\\nocite{}) on 2002, Jan.~29.40 UT (see Nakano 2002)\\nocite{n02}.  Mazzali et al. (2002) inferred an explosion date of 2002, Jan.~$28\\pm 0.5$ UT.  Attracted by the broad spectral features (e.g. Kinugasa et al.  2002; Meikle et al. 2002)\\nocite{kka+02,mls+02} we began observing the SN at the Very Large Array (VLA\\footnotemark\\footnotetext{The VLA is operated by the National Radio Astronomy Observatory, a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}). ",
        "conclusions": "\\label{sec:conc}  The type Ic SN\\,1998bw, likely associated with GRB\\,980425, was peculiar in two ways. It  exhibited broad photospheric absorption lines, and exceedingly strong radio emission at early times.  The latter was interpreted to arise from relativistic ($\\Gamma\\sim$ few) ejecta of total energy $E_\\Gamma\\age 10^{49}$ erg. These two peculiarities made sense in that the simple theory suggested that broad photospheric features are a reliable indicator of relativistic ejecta, a necessary condition for $\\gamma$-ray emission.  The type Ic SN\\,2002ap elicited much interest because it too displayed similar broad lines.  However, from our radio observations we estimate the energy in relativistic electrons and magnetic fields to be quite modest: $E\\approx 1.5\\times 10^{45}$ ergs in ejecta with a velocity $\\approx 0.3c$. In view of this, the absence of $\\gamma$-rays from \\sn{} is not surprising (Hurley et al. 2002)\\nocite{hur02}.  Both the energy and speed of the ejecta can be accounted for in the standard hydrodynamical model.  We thus conclude that broad photospheric lines are not good predictors of relativistic ejecta.  Separately, the broad photospheric features led modelers to conclude that \\sn{} was a hypernova with an explosion energy of $E_{51}\\sim 4-10$ erg (Mazzali et al. 2002). However, the radio observations paint a different picture.  The low $E_\\Gamma$ make \\sn{} an ordinary Type Ib/c SN (or perhaps even a low energy event; \\S\\ref{sec:general}).  We also draw attention to the significant role played by asymmetries in forming the optical lines which, when not properly accounted for can lead to discrepant estimates (cf.~the large variations in the estimated SN energy release for SN\\,1998bw; see H\\\"oflich, Wheeler \\& Wang 1999\\nocite{hww99}; Iwamoto et al. 1998\\nocite{imn+98}).  Thus, we wonder how reliable are the inferred $E_{51}$ and $M_1$ values.  Along the same vein, we note that Kawabata et al. (2002)\\nocite{kji+02} suggest, based on spectro-polarimetric observations, a jet with a speed of $0.23c$ and carrying $2\\times 10^{51}$ erg.  Such a jet, regardless of geometry, would have produced copious radio emission.  We end with three conclusions.  First, at least from the perspective of relativisic ejecta, \\sn{} was an ordinary Ib/c SN.  Second, broad photospheric lines appear not to be a good proxy for either an hypernova origin or $\\gamma$-ray emission.  Third, radio observations offer a practical and accurate proxy for relativistic ejecta."
    },
    "0206/astro-ph0206460_arXiv.txt": {
        "abstract": "We present model lightcurves which have been created in order to explain the orbital modulation observed in the radio emission of Cyg X-1. We invoke variable absorption by the stellar wind as the black hole jet orbits around the OB companion star and find that a very simple model is able to reproduce the amplitudes and frequency dependence of the observed lightcurves. ",
        "introduction": "Cygnus X-1 was first observed to be a radio source by Braes \\& Miley (1971). At the time of discovery the radio flux was $\\sim 20$ mJy and it generally maintains a constant mean level of $\\sim$ 10--25 mJy at cm wavelengths while the system is in its usual low/hard X-ray spectral state. However, when Cyg X-1 makes one of its rare state changes to the soft state the radio emission appears to be quenched and drops to below a detectable level (Pooley 2001; Brocksopp et al. 1999a, and references within). This behaviour is similar to that of GX 339$-$4 -- another black hole X-ray binary which has been monitored in the radio throughout a soft state period and transitions (Fender et al. 1999, Corbel et al. 2000).  An important feature of the radio emission is the observed flat synchrotron spectrum to millimetre wavelengths (Fender et al. 2000) with no evidence for a cut-off at higher frequencies. Unfortunately the infrared emission is dominated by the supergiant and so it is not possible to determine how far the flat spectrum continues. This spectrum suggests that Cyg X-1 has a quasi-continuous jet of the type modelled by Hjellming \\& Johnston (1988) and/or Blandford \\& K\\\"onigl (1979). These predictions have recently been confirmed with the VLBI imaging of the jets of Cyg X-1 (Stirling et al. 2001).  Two types of periodic behaviour can be observed in the radio emission. The longer of these is at $\\sim$140 days (although is probably not stable in the long term) and is possibly due to precession of the disc-jet system (Brocksopp et al. 1999a); alternatively a variable mass accretion rate such as that proposed for LMC X-3 (Wilms et al. 2001; Brocksopp, Groot \\& Wilms 2001) is also a potential cause of this long period. The 5.6-day orbital period can also be detected in the radio emission (Pooley, Fender \\& Brocksopp 1999) and it is this modulation that we now investigate.  The orbital period of Cyg X-1 has been well-established for more than two decades and particularly since the work of Gies \\& Bolton (1982). The ephemeris has been revised recently (e.g. Brocksopp et al. 1999b, Sowers et al. 1998, LaSala et al. 1998) but the results have all been consistent with each other.  X-ray data have also revealed the orbital period (Paciesas et al. 1997) and in Brocksopp et al. (1999a) we discussed the significance of finding the orbital period at all wavelengths in the Cyg X-1 system. Detection of the orbital period of black hole X-ray binaries in optical and infrared photometry is standard and double-peaked orbital lightcurves, due to the gravitational pull on the star by the black hole, are expected -- minima occur in the lightcurves at the two conjunctions. However, it is not so common to detect the orbital period in X-ray or radio data, particularly since Cyg X-1 is a non-eclipsing system and is thought to have a circular orbit. Thus Brocksopp et al. (1999a) postulated that these modulations could be the result of absorption by the stellar wind. This was later confirmed by Wen et al. (1999) who studied the orbital modulation of the $RXTE$ soft X-ray data and produced a relatively successful model of wind absorption. The X-ray orbital lightcurves have also been studied by Ba{\\l}uci\\'{n}ska--Church et al. (2000) who found that the duration and phase of the minima are also variable and likely to be due to `clumpiness' in the stellar wind.  At radio wavelengths we found that the modulation is frequency-dependent with the strongest modulation at higher frequencies -- contrary to what one would expect for straightforward free-free absorption of a point source. There is also a phase lag present, the duration of which may also be frequency-dependent. These observations add support to the proposed continuous jet model on account of an extended jet-like structure probing different densities of an absorbing medium. Furthermore it would be unlikely that the modulations could be maintained as consistently as observed, were the radio emission produced in discrete ejections.  We now produce a model of wind absorption that can explain the observed radio emission at three wavelengths. To do this we use the focussed stellar wind model of Friend \\& Castor (1982), which they and Gies \\& Bolton (1986b) have shown to be applicable to the Cyg X-1 system. In this model the spherical stellar wind of the supergiant becomes distorted by the gravitational field of the black hole, thus resulting in enhanced emission occuring close to the axis joining the supergiant and black hole. ",
        "conclusions": "We have presented a model to explain the orbital modulation observed in the radio emission of Cyg X-1. By invoking absorption of the radio emission by the stellar wind of the companion star we are able to produce model lightcurves with amplitudes and frequency dependence comparable with the real data. We acknowledge that the fits could be improved with the use of a more sophisticated stellar wind model -- the production of such a wind model is work in progress."
    },
    "0206/astro-ph0206197_arXiv.txt": {
        "abstract": "{As part of our program to map the large-scale distribution of galaxies behind the Milky Way, we used the Parkes 210\\,ft (64\\,m) radio telescope for pointed \\HI\\ observations of a sample of low surface-brightness (due to heavy obscuration) spiral galaxies selected from the deep optical Zone of Avoidance (ZOA) galaxy catalog in the Hydra/Antlia region (Kraan-Korteweg 2000a). Searching a simultaneous velocity range of either 300 to 5500\\kms\\ or 300 to 10\\,500\\kms\\ to an rms level of typically $2 - 4$\\,m\\,Jy resulted in detections in 61 of the 139 pointings, leading to a total of 66 detections (an additional detection was made in a reference position, and two other pointings revealed two and four independent signals respectively). Except for 2 strong \\HI\\ emitters identified in the shallow Zone of Avoidance \\HI\\ survey (Henning et al. 2000), all \\HI\\ detections are new. An analysis of the properties of the observed and detected galaxies prove that pointed \\HI\\ observations of highly obscured galaxies allow the tracing of a population of nearby, intrinsically large and bright spiral galaxies that otherwise would not be recovered. The new data identified a previously unrecognized nearby group at $\\ell \\sim 287\\fdg5, b \\sim-9\\fdg5, V\\sim 1700$\\kms, the continuation of the Hydra/Antlia filament on the opposite side of the Galactic plane, and helped to delimit a distinct void in the ZOA centered at 2000\\kms. ",
        "introduction": "Until recent years, galaxy surveyors have avoided the difficult territory of the Milky Way.  Great strides toward filling in the low-galactic-latitude gap in our knowledge of the galaxian distribution have been made by optical, radio, near- and far-infrared, and X-ray surveys (see review by Kraan-Korteweg \\& Lahav 2000, and the many contributions in ``Mapping the Hidden Universe'', ASP Conf. Ser. 218, eds. Kraan-Korteweg \\etal 2000).  The first step in this cartography (with the exception of 21-cm surveys which immediately yield the angular coordinates and redshift of galaxies) is the two-dimensional mapping of galaxies.  Follow-up redshift measurements are required to map the distribution of galaxies in three dimensions.  The two-dimensional galaxy distribution alone can be misleading.  For instance, the prominent overdensity of galaxies in Vela ($\\ell \\sim 280\\deg, b \\sim +6\\deg$) made apparent by the optical search by Kraan-Korteweg (2000a) in the Hydra/Antlia extension was found to be due to a superposition of a nearby ($\\sim 3000$\\kms) filament connecting to the Hydra cluster, a more distant ($\\sim 6000$\\kms) shallow extended supercluster, and a very distant ($\\sim 16\\,000$\\kms) wall-like structure crossing the ZOA (Kraan-Korteweg \\etal 1995).  In this paper, we will present the results from pointed \\HI\\ observations of a sample of low surface-brightness (LSB) obscured spiral galaxies selected from the deep optical galaxy catalog in the ZOA ($266\\deg \\la \\ell \\la 296\\deg, -10\\deg \\la b \\la +8\\deg$) in the extension of the Hydra and Antlia clusters (Kraan-Korteweg 2000a). The optical catalog is the first of a series of five systematic galaxy searches covering the southern Milky Way between Galactic longitudes $245\\deg$ and $350\\deg$ (see Fig.~1 in Kraan-Korteweg 2000a). Galaxies were identified by eye through the systematic inspection of 50-times amplified images of the IIIaJ film copies of the ESO/SRC sky survey. Although the majority of the newly uncovered galaxies are small and faint, many of them may be intrinsically large and nearby but reduced in apparent size, magnitude and classifiable morphology due to the increasing foreground dust absorption at lower latitudes.  An analysis of the catalogs (see Kraan-Korteweg 2000a and Woudt \\& Kraan-Korteweg 2001 for details) show that galaxies above a diameter limit of $D \\ga 0\\farcm2$ (determined at an isophote of approximately 24.5 mag/arcsec$^2$) can be discovered ``easily'' through obscuration layers of 3 magnitudes of extinction. Up to that extinction level of $A_{\\rm B} = 3\\fm0$, the catalogs are complete for galaxies with intrinsic diameters of $D^{\\rm o} \\ge 1.3$~arcmin (extinction-corrected, i.e. the diameter they would have if they were not lying behind the Milky Way). The optical searches succeed in a reduction of the solid angle of the ZOA by a factor of about $2 - 2.5$, respectively from $A_{\\rm B} = 1\\fm0$ to $A_{\\rm B} = 3\\fm0$ (see also Fig.~\\ref{dist}).  In order to also reduce the gap in redshift space, redshifts are required for a representative sample of this newly optically filled-in part of the ZOA. For the determination of the peculiar velocity of the Local Group and the mapping of the velocity flow field, the nearby galaxy population is particularly important. We obtain individual optical spectroscopy of the brightest galaxies with the SAAO 1.9\\,m telescope (Kraan-Korteweg \\etal 1995 for the Hydra/Antlia ZOA region; Fairall \\etal 1998, and Woudt \\etal 1999 for the Crux and Great Attractor regions respectively) and 21-cm line observations of extended LSB spirals with the 64m Parkes radio telescope (this paper), and finally, low resolution, multifiber spectroscopy for the high-density regions (Felenbok et al.~1997, Woudt \\etal 2002). We typically measure recession velocities for about 15\\% of our ZOA galaxies.  These three observing methods are complementary in galaxy populations, characteristic magnitude, diameter range, and the depth of volume they probe: whereas the multifiber spectroscopy gives a good description of clusters and dense groups in the ZOA out to recession velocities of 25\\,000\\,\\kms, the SAAO and \\HI\\ observations cover the bright end of the galaxy distribution and provide a fairly homogeneous sampling of galaxies out to 10\\,000\\,\\kms (Kraan-Korteweg \\etal 1994).  The \\HI\\ observations are vital in recovering an important fraction of the nearby spiral galaxy population which would otherwise be impossible to map.  In addition to the extremely obscured and/or LSB galaxies, some of the brighter spiral galaxies with redshifts known from optical spectroscopy were reobserved with the Parkes radiotelescope, as their \\HI\\ data are relevant to our program of mapping the peculiar velocity field in the ZOA via the Tully--Fisher relation.  Our pointed \\HI\\ observations are complementary to the ongoing blind \\HI\\ survey in the ZOA, conducted with the multibeam receiver on the Parkes telescope (Henning \\etal 2000, Staveley-Smith \\etal 2000b). Unlike the current work, the blind \\HI\\ ZOA survey is not optically selected, but is rather a fully-sampled survey over the region $212\\deg \\le \\ell \\le 36\\deg$, $|b| \\leq 5\\deg$.  The velocity coverage of the blind survey is $-1200$ to 12\\,700\\,\\kms, and the sensitivities obtained by the deep blind survey will be similar to those obtained in the current work.  Note that most of our objects lie outside of the blind \\HI\\ survey region.  In the following section, a description of the observations is given. Sect.~3 then provides the \\HI\\ data and profiles of the detected galaxies, including details about a number of pointings which revealed more than one signal, two of which do not seem to have an optical counterpart even though the extinction is not extreme at those positions. In Sect.~4 the observed galaxies that were not detected in \\HI\\ are listed with the searched velocity range. This is followed by an analysis of the properties of the detected (and non-detected) galaxies (Sect.~5) and a brief discussion of the resulting galaxy distribution in redshift space in relation to known features adjacent to the ZOA (Sect.~6), and finally, a summary is provided in Sect.~7. ",
        "conclusions": "The pointed 21-cm line observations of partially- to heavily-obscured spiral galaxies uncovered in the deep optical search for galaxies in the ZOA has proven this technique to be very powerful for mapping a population of galaxies -- which on average actually are intrinsically large and bright -- whose redshifts would be difficult or impossible to obtain otherwise.  Besides the chance detection of a few galaxies without optical counterparts (see Sect.~\\ref{weirdt}), the 66 detections -- and the distribution of the non-detections -- have helped to delimit a number of suspected or previously unknown large-scale structures, in particular the continuation of the Hydra/Antlia bridge across the GP and the void with a radius of about 1000\\kms\\ centered in the ZOA at about 2000\\kms.  These HI-observations helped reduce the width of the ZOA in redshift space. They will also be invaluable in mapping the peculiar velocity field in the ZOA. The missing gap in LSS ($|b| \\la 5\\deg$) will be filled with the detections from the deep ZOA survey that is being carried out, at similar sensitivities, with the multibeam receiver on the Parkes radiotelescope."
    },
    "0206/astro-ph0206409.txt": {
        "abstract": "{We present a Monte-Carlo package for simulation of high-redshift supernova data, SNOC. Optical and near-infrared photons from supernovae are ray-traced over cosmological distances from the simulated host galaxy to the observer at Earth. The distances to the sources are calculated from user provided cosmological parameters in a Friedmann-Lema\\^{\\i}tre universe, allowing for arbitrary forms of ``dark energy''.  The code takes into account gravitational interactions (lensing) and extinction by dust, both in the host galaxy and in the line-of-sight. The user can also choose to include exotic effects like a hypothetical attenuation due to photon-axion oscillations.  SNOC is primarily useful for estimations of cosmological parameter uncertainties from studies of apparent brightness of Type Ia supernovae vs redshift, with special emphasis on potential systematic effects.  It can also be used to compute standard cosmological quantities like luminosity distance, lookback time and age of the universe in any Friedmann-Lema\\^{\\i}tre model with or without quintessence. ",
        "introduction": "The study of brightness of Type Ia supernovae (SNe) at high redshifts has become one of the most important tools in observational cosmology over the last few years \\citep{GoobarPerlmutter1995,Perlmutter1999,Riess1998,Riess2001}, giving the first direct observational evidence for a presently accelerating universe. With the increased capabilities of observing very high-redshift SNe, more comprehensive analysis tools are needed \\citep[see, e.g.,][]{sn97ff}.  In this work we describe a Monte-Carlo simulation package called {\\em The SuperNova Observation Calculator} (SNOC) which produces synthetic samples of SN observations that can be used to estimate the accuracy of the magnitude-redshift method for measuring cosmological parameters as well as quantifying the possibility of constraining extragalactic dust properties \\citep{dust} or investigating the matter distribution in the Universe \\citep{compact}.  For example, SNOC can be used to quantify the deviations in the Hubble diagram of Type Ia SNe due to gravitational lensing or dust extinction along the line-of-sight. Tools are also provided to compute the possible contamination of an observational sample by core collapse SNe and the likelihood of observation of multiple (lensed) images from individual high-$z$ SNe \\citep{Goobar:2002}. ",
        "conclusions": "As we enter a phase of high-precision cosmological measurements, sophisticated tools for data analysis are required, especially to evaluate potential systematic effects related to the method. The Monte-Carlo simulation package SNOC is mainly designed for the study of the precision and possible biases of measured cosmological parameters from high-$z$ Type Ia SNe. Besides calculating the luminosity distance for arbitrary cosmological parameters: $H_0, \\Omega_M, \\Omega_X, \\alpha_X(z)$, the code also allows for the estimates of observed magnitudes for high-$z$ SNe taking into account the inhomogeneity of matter and dimming by, e.g., dust or more exotic process such as photon-axion oscillations.  Along with the simulation, a maximum-likelihood analysis package has been developed for fitting of cosmological parameters and possible bias.  %For information on how to download the code, please contact the %authors."
    },
    "0206/astro-ph0206368_arXiv.txt": {
        "abstract": "The early-time optical spectrum of the Type Ic supernova (SN) 2002ap was characterized by unusually broad features, leading some authors to designate it a ``hypernova.\" We present optical spectropolarimetry of this object 16 and 37 days after the estimated date of explosion.  After correcting for interstellar polarization, we find evidence for a high level of intrinsic continuum polarization at both epochs: $p \\gtrsim 1.3\\%$ on day 16 and $p \\gtrsim 1.0\\%$ on day 37. Prominent line polarization is also seen, especially in the trough of the Ca~II near-infrared triplet during the second epoch. When interpreted in terms of the oblate, electron-scattering model atmospheres of H\\\"{o}flich (1991), our results imply an asymmetry of at least $20\\%$ (day 16) and $17\\%$ (day 37).  The data suggest a fair degree of axisymmetry, although the polarization angle of the dominant polarization axes are different by about $55^\\circ$ between the two epochs, implying a complex morphology for the thinning ejecta.  In particular, there exists some spectropolarimetric evidence for a different distribution of Ca relative to iron-group elements.  We also present flux spectra of SN~2002ap taken 131 and 140 days after the explosion.  The spectra are characterized by a very weak continuum and broad emission lines, indicating that SN~2002ap has entered the nebular phase.  The spectral features are substantially similar to those of ``normal'' SNe Ic in the nebular phase, and the emission lines are not significantly broader. However, some of the broad lines are very sharply peaked, and may possess a narrow component (probably unresolved by our spectra, FWHM $\\lesssim 400$ \\kms) that is redshifted by $\\sim 580$ \\kms\\ with respect to the systemic velocity of the host galaxy. ",
        "introduction": "\\label{sec:introduction}  Although core-collapse supernovae (SNe) present a wide range of spectral and photometric properties, there is growing consensus that much of this variety is due to the state of the progenitor star's hydrogen and helium envelopes at the time of explosion. Those stars with massive, intact envelopes produce SNe of Type II-plateau (SNe II-P; II due to the presence of H in the spectrum, and ``plateau'' due to the characteristic shape of the light curve), those that have lost their entire hydrogen envelope (perhaps through stellar winds or mass transfer to a companion) result in SNe~Ib (Type I because of the absence of obvious hydrogen in their spectra), and those that have been stripped of both hydrogen and most (or all) of their helium produce SNe~Ic; see Filippenko (1997) for a general review.  Recently, a new subclass of objects has emerged whose members generically resemble SNe~Ic (no hydrogen or obvious helium spectral features), but, unlike traditional SNe~Ic, have spectra characterized by unusually broad features at early times, indicating velocities in excess of $\\sim 30,000$ \\kms.  A few also possess inferred kinetic energies exceeding that of ``normal'' core-collapse SNe by more than a factor of 10 (see, e.g., Nomoto et al. 2001).  These objects are colloquially referred to as ``hypernovae,'' although not all of them are clearly more luminous or energetic than normal SNe~Ic.  There are currently 5 generally accepted members of this rare class: SN~1997dq, SN~1997ef, SN~1998bw, and, most recently, SN 2002ap and SN~2002bl (see, e.g., Iwamoto et al. 2000; Matheson et al. 2000b; Filippenko, Leonard, \\& Moran 2002; Mazzali et al. 2002; Kawabata et al. 2002).  A related subclass of SNe exhibits many of the characteristics of these objects, but with hydrogen present in the spectra; the clearest examples are SN 1997cy and SN 1999E (see Filippenko 2001, and references therein), and they, too, are sometimes called hypernovae.  Intense interest in hypernovae has been sparked not only by their peculiar spectral features, but also by the strong spatial and temporal association between the brightest and most energetic of these events, SN 1998bw, and the $\\gamma$-ray burst (GRB) 980425 (e.g., Galama et al. 1998).  This potential association has fueled the proposition that some (or, perhaps all) core-collapse SNe explode due to the action of a ``bipolar'' jet of material (Wheeler, Meier, \\& Wilson 2002; Khokhlov et al. 1999; MacFadyen \\& Woosley 1999), as opposed to the conventional neutrino-driven mechanism (Colgate \\& White 1966; Burrows et al. 2000, and references therein).  Under this paradigm, a GRB is only produced by those few events in which the progenitor has lost all of its outer envelope material (i.e., it is a ``bare core'' collapsing), and is only observed if the jet is closely aligned with our line-of-sight (l-o-s). With such an explosion mechanism, one can well imagine that severe distortions from spherical symmetry may exist in the ejecta.  Since all hot, young SN atmospheres are dominated by electron scattering (Wagoner 1981), which is an inherently polarizing process, spectropolarimetry provides a powerful probe of young SN morphology.  The basic question is whether the SN is round: by symmetry, if the atmosphere is spherical the directional components of the electric vectors cancel to produce zero net polarization.  In contrast, an asymmetric atmosphere will yield a nonzero polarization due to incomplete cancellation.  From recent polarimetry of core-collapse events, a tentative trend has been identified: the degree of polarization (and, hence, asymmetry) increases with decreasing progenitor envelope mass (Wheeler 2000). Indeed, while the intrinsic polarizations of ``normal'' SNe II-P are generally below $1\\%$ (Leonard et al. 2001; Leonard \\& Filippenko 2001), a polarization of over 4\\% has been reported for an SN Ic, suggesting an axial asymmetry of more than 3:1 in this event (Wang et al. 2001).  The number of events investigated in detail with spectropolarimetry, however, remains small.  An additional technique that has been used to infer explosion asymmetry is the analysis of nebular line profiles.  Recent simulations involving explosive nucleosynthesis in aspherical, jet-induced SN explosions predict that intermediate-mass and heavy elements such as iron are ejected (at high velocity) primarily along the poles whereas elements synthesized in the progenitor (e.g., He, C, Ca, O) are preferentially located at lower velocities near the equatorial plane in the expanding ejecta (Maeda et al. 2002; Khokhlov \\& H\\\"{o}flich 2001; H\\\"{o}flich et al. 2001).  Maeda et al. (2002) model the effect that explosion asymmetry has on the spectra of hypernova explosions, and compare the results to the observed nebular line profiles of SN~1998bw.  The main observable diagnostic for explosion asymmetry is found to be the ratio of the width of a probable [\\ion{Fe}{2}] blend near 5200~\\AA\\ to that of [\\ion{O}{1}] $\\lambda\\lambda$6300, 6364~\\AA.  Maeda et al. (2002) find models of spherical explosions to be incapable of yielding velocity ratios greater than $\\sim$ 3:2, whereas aspherical models can generate ratios of more than 2:1.  Although a specific degree of asphericity is not given, Maeda et al. (2002) find the nebular-phase line-width ratio of [\\ion{Fe}{2}] to [\\ion{O}{1}] for SN~1998bw to be inconsistent with a spherical explosion.  In this paper, we report on spectropolarimetric observations obtained for the recent peculiar\\footnote{Berger, Kulkarni, \\& Chevalier (2002) recently used radio observations to conclude that SN~2002ap does not exhibit evidence for a large amount of relativistic ejecta, and hence is an ``ordinary\" SN~Ic; however, we feel that it remains a somewhat peculiar SN~Ic, given its unusually broad early-time spectral features.}  Type Ic SN~2002ap in M74 (Kinugasa et al. 2002; Filippenko \\& Chornock 2002; Mazzali et al.  2002; Gal-Yam, Ofek, \\& Shemmer 2002) during the photospheric phase, and two nebular-phase optical flux spectra.  We describe the observations in \\S~\\ref{sec:observations}, present our results and analysis in \\S~\\ref{sec:resultsandanalysis}, and discuss our conclusions in \\S~\\ref{sec:conclusions}.  Additional details of the observations and analysis are given in Appendix~A.  As this paper was nearing completion, a submitted paper by Kawabata et al. (2002) describing similar spectropolarimetric data appeared on astro-ph; we briefly compare our results with that work in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  There is strong evidence from spectropolarimetry and tentative evidence from nebular-phase line profiles that SN~2002ap is not spherical.  It is useful to compare the spectropolarimetry of SN~2002ap with the small but growing number of core-collapse events for which polarimetry measurements exist.  The lower bounds on the continuum polarization level of SN~2002ap ($1.3\\%$ and $1.0\\%$ on days 16 and 37, respectively) are larger than those typically derived for SNe II-P (Wheeler 2000), but not nearly as great as that inferred for the SN~Ic 1997X (Wang et al. 2001).\\footnote{The presentation of data for SN~1997X by Wang et al. (2001) shows broadband averages of the polarization data.  It is therefore unknown what fraction of the claimed polarization ($p > 4\\%$) might be due to line-feature modulations similar to what is seen in the \\ion{Ca}{2} IR triplet of SN 2002ap during the second epoch, and not to intrinsically polarized continuum light.}  A more direct comparison to a similar event can be made with the spectropolarimetry presented by Patat et al. (2001) for the hypernova SN~1998bw.  Patat et al. (2001) present spectropolarimetry from two epochs, $-7$ and +10 days after $B$-band maximum (our observations of SN~2002ap occurred on days 7 and 28 relative to the date of $B$ maximum derived by Gal-Yam et al. 2002) with coverage from 4000~\\AA\\ to 6900~\\AA.  Although the ISP is unknown for this event,\\footnote{The Galactic ISP derived by adopting the measured linear polarization of a nearby star, HD 184100 (Kay et al. 1998), is probably not appropriate since this star has a spectroscopic parallax of only $\\sim 14$ pc, which is insufficient to fully sample the dust in the Galactic plane.} we note that the overall level of the observed polarization ($p \\approx 0.5 - 0.7\\%$), and the marked increase in the polarization level from blue to red wavelengths (particularly in the earliest epoch), are quite similar to the behavior of SN~2002ap.  In addition to the high inferred continuum polarization, another interesting finding for SN~2002ap is that while the data from both epochs suggest a substantial degree of axisymmetry (Fig.~\\ref{fig:3}), the dominant axes have quite different polarization angles, changing from $\\sim 95^\\circ$ on day 16 to $\\sim 150^\\circ$ on day 37 (\\S~\\ref{sec:removalofisp}).  This behavior is different from that observed in other core-collapse events, most notably those resulting from progenitors with massive envelopes intact at the time of explosion.  For SN~1999em, a classic Type II-P event, Leonard et al. (2001) show that a common, wavelength-independent P.A. exists for all observational epochs, which includes days 12, 45, 54, 164, and 168 after explosion.  This suggests a more complex morphology for the thinning ejecta of SN~2002ap.  As mentioned in \\S~\\ref{sec:introduction}, a potentially relevant result of recent simulations involving explosive nucleosynthesis in aspherical, jet-induced SN explosions is that intermediate-mass and heavy elements (e.g., Fe) as well as freshly synthesized radioactive material such as $^{56}$Ni are ejected (at high velocity) primarily along the poles whereas elements synthesized in the progenitor (e.g., He, C, Ca, O) are preferentially located at lower velocities near the equatorial plane in the expanding ejecta (Maeda et al. 2002; Khokhlov \\& H\\\"{o}flich 2001; H\\\"{o}flich et al. 2001).  Although the P.A. in the ISP-corrected data from day 37 (bottom, right panel of Fig.~\\ref{fig:4}) is difficult to define in some regions due to the low polarization level, there is evidently a change near $\\lambda \\approx 6000$ \\AA.  Below this wavelength, $\\theta \\approx 100^\\circ$, a value similar to the one that characterizes the day 16 data.  Above $\\sim 6000$ \\AA, $\\theta \\approx 150^\\circ$ tends to dominate; this is especially easy to see (and believe, due to the high polarization) in the \\ion{Ca}{2} IR absorption trough (see also Fig.~\\ref{fig:3}).  Might this change in P.A. be due to the different spatial distributions of the iron-group elements, which dominate at blue wavelengths, from the elements synthesized by the progenitors, such as Ca and O, that are responsible for much of the opacity at red wavelengths (see Fig.~\\ref{fig:2}a)? Along these same lines, an explanation for the nearly constant P.A. on day 16 at all wavelengths, and its similarity to the P.A. of day 37 at blue wavelengths, could be due to the photosphere at early times being located in the highest-velocity (i.e., polar) material. By day 37, perhaps the SN has cooled sufficiently to expose the inner, slower-moving equatorial regions, and we are witnessing the competition between these two different distributions of material.  To be sure, a small change in the ISP estimate, or even the wavelength of maximum polarization of the ISP, could alter the inferred P.A. of the continuum in the second epoch considerably.   Further interpretation of the polarization change near $6000$ \\AA\\ in the second epoch is probably not warranted.  The $\\sim 55^\\circ$ P.A. difference between the the day 16 data and the region associated with the \\ion{Ca}{2} IR trough on day 37, however, remains regardless of the ISP choice, and may well indicate a true difference in the distribution of Ca relative to iron-group elements.   \\begin{figure} \\ssp \\begin{center} \\rotatebox{0}{ \\scalebox{0.7}{ \\plotone{leonard.fig8.ps} } } \\end{center} \\caption{The polarized flux of SN~2002ap on 2002 Feb. 14 (16 days after explosion; {\\it thin line}) compared with the total flux spectrum from the same date redshifted by 0.23$c$ ({\\it thick line}), arbitrarily scaled and offset for comparison of features. \\label{fig:8} } \\end{figure}  Just prior to completion of this research, a paper describing similar data appeared on astro-ph by Kawabata et al. (2002).  Although the observational epochs are slightly different (Kawabata et al. present data from 2002 February 9.2--11.3 and March 8.2--10.2), the general results are similar: an observed polarization level in the February epochs of $p \\approx 0.6\\%$ that may rise somewhat toward the red (the increase is not as obvious as it is in our data since their data does not extend beyond 8300~\\AA) and a dramatic polarization increase in the \\ion{Ca}{2} IR triplet's absorption trough followed by depolarization in the emission profile in the March data.  One interesting speculation made by Kawabata et al. (2002) is that much of the polarized continuum at early times may result from scattering off of electrons in a jet, or bipolar jets, of material emitted from the SN during the explosion. Supporting this hypothesis is the general similarity demonstrated between the intrinsic polarized flux (i.e., $p_{\\rm intrinsic} \\times f$) and the total flux spectrum redshifted by $0.23c$, the presumed speed of the jet.  We also find qualitative agreement between the intrinsic polarized flux and the total flux redshifted by $z = 0.23c$ during our February epoch (Fig.~\\ref{fig:8}), although we have not investigated how significant the correlation is.  We do note, though, that the speculation by Kawabata et al. (2002) is consistent with the scenario proposed above that the polarization P.A. on day 16 is dominated by the distribution of the polar jet material.  Countering the jet hypothesis, however, is the recent assertion by Berger et al. (2002) that the extreme faintness of SN~2002ap at radio wavelengths is inconsistent with the presence of a jet, regardless of the viewing geometry.  In any event, additional spectropolarimetry of young SNe Ic is certainly warranted in order to further test the jet model.  The total flux spectra obtained in the early-nebular phase confirm that SN~2002ap continues the rapid spectral evolution previously witnessed during the photospheric phase by Mazzali et al. (2002).  It is characterized by a very weak continuum and broad emission lines, which demonstrate that it is already well along the transition to the fully nebular phase.  The lines have widths similar to those observed in ``normal'' nebular SNe Ic.  There are, however, unusual narrow lines superimposed on some of the broad-line profiles, including especially those of [\\ion{O}{1}] $\\lambda\\lambda$6300, 6364 and \\ion{Mg}{1}] $\\lambda$4571.  It is possible that the ratio of the width of the [\\ion{Fe}{2}] blend near 5200~\\AA\\ compared with that of [\\ion{O}{1}] $\\lambda\\lambda$6300, 6364 is indicative of asymmetry in the line-emitting region, but additional, later-time spectra are needed to confirm this result."
    },
    "0206/astro-ph0206018_arXiv.txt": {
        "abstract": "The results of {\\chandra} snapshot observations of 11 LINERs (Low-Ionization Nuclear Emission-line Regions), three low-luminosity Seyfert galaxies, and one HII-LINER transition object are presented. Our sample consists of all the objects with a flat or inverted spectrum compact radio core in the VLA survey of 48 low-luminosity AGN (LLAGN) by Nagar et al. (2000). An X-ray nucleus is detected in all galaxies except one and their X-ray luminosities are in the range $5\\times10^{38}$ to $8\\times10^{41}$ {\\eps}. The X-ray to H$\\alpha$ luminosity ratios for 11 out of 14 objects are in good agreement with the value characteristic of LLAGNs and more luminous AGNs, and indicate that their optical emission lines are predominantly powered by a LLAGN. For three objects, this ratio is less than expected.  Comparing with multi-wavelength results, we find that these three galaxies are most likely to be heavily obscured AGN. We compare the radio to X-ray luminosity ratio of LLAGNs with those of more-luminous AGNs, and confirm the suggestion that a large fraction of LLAGNs are radio loud.  \\textbf{} ",
        "introduction": "Low-Ionization nuclear emission-line regions (LINERs) are found in many nearby bright galaxies (e.g., Ho, Filippenko, \\& Sargent 1997a). Extensive studies in various wavelengths have shown that type 1 LINERs (LINER 1s, i.e., those galaxies having broad H$\\alpha$ and possibly other broad Balmer lines in their nuclear optical spectra) are powered by a low-luminosity AGN (LLAGN) with a bolometric luminosity less than $\\sim10^{42}$ {\\eps} (Ho et al. 2001; Terashima, Ho, \\& Ptak 2000a; Ho et al. 1997b). On the other hand, the energy source of LINER 2s is likely to be heterogeneous.  Some LINER 2s show clear signatures of the presence of an AGN, while others are most probably powered by stellar processes, and the luminosity ratio {\\LX}/{\\LHa} can be used to discriminate between these power sources (e.g., Terashima et al. 2000b).  It is interesting to note that currently there are only a few LINER 2s known to host an obscured AGN (e.g., Turner et al. 2001). This paucity of obscured AGN in LINERs may indicate that LINER 2s are not simply a low-luminosity extension of luminous Seyfert 2s, which often show heavy obscuration with a column density averaging {\\NH} $\\sim$ $10^{23}$ {\\pcm} (e.g., Turner et al. 1997). Alternatively, biases against finding heavily obscured LLAGNs may be important. For example, objects selected through optical emission lines or X-ray fluxes are probably biased in favor of less absorbed ones, even if one uses the X-ray band above 2 keV.  In contrast, radio observations, particularly at high frequency, are much less affected by absorption. Nagar et al. (2002) have reported a VLA 2 cm radio survey of all 96 LLAGNs within a distance of 19 Mpc. These LLAGNs come from the Palomar spectroscopic survey of bright galaxies (Ho et al. 1997a). As a pilot study of the X-ray properties of LLAGNs, we report here a {\\chandra} survey of a subset, comprising 14 galaxies, of Nagar et al's (2002) sample. We have detected 13 of the galactic nuclei with {\\chandra}. We also examine the ``radio loudness'' of our sample and compare it with other classes of AGN. ",
        "conclusions": "\\subsection{Power Source of LINERs}  We test whether the detected X-ray sources are the power source of their optical emission lines by examining the luminosity ratio {\\LX}/{\\LHa}. The H$\\alpha$ luminosities ({\\LHa}) were taken from Ho et al. (1997a) and corrected for the reddening estimated from the Balmer decrement for the narrow lines. The X-ray luminosities ({\\LX}) in the $2-10$ keV band, and corrected for absorption, were used.  The resulting {\\LX}/{\\LHa} ratios of most objects are in the range of AGNs ($\\log$ {\\LX}/{\\LHa} {\\simgt}1) and in good agreement with the strong correlation between {\\LX} and {\\LHa} for LLAGNs, luminous Seyferts, and QSOs presented in Terashima et al. (2000a) and Ho et al. (2001). This indicates that their optical emission lines are predominantly powered by a LLAGN.  The three objects NGC 2787, NGC 5866, and NGC 6500, however, have much lower {\\LX}/{\\LHa} ratios ($\\log$ {\\LX}/{\\LHa} $\\simlt$0) than expected from the correlation, and their X-ray luminosities are not enough to power the H$\\alpha$ luminosities. This X-ray faintness could indicate one or more of several possibilities such as (1) an AGN is the power source, but is heavily absorbed at energies above 2 keV, (2) an AGN is the power source, but is currently switched-off or in a faint state, and (3) the optical narrow emission lines are powered by some other source(s) than an AGN.   If an AGN is present in these X-ray faint objects and absorbed in the hard energy band above 2 keV, only scattered and/or highly absorbed X-rays can be observed, and then the intrinsic luminosity would be much higher than that observed. This can account for the low {\\LX}/{\\LHa} ratios and high radio to X-ray luminosity ratios ($\\nu L_{\\nu}$(5 GHz)/{\\LX}). If the intrinsic X-ray luminosities are about one or two orders of magnitude higher than those observed, as is often inferred for Seyfert 2 galaxies, {\\LX}/{\\LHa} and $\\nu L_{\\nu}$(5 GHz)/{\\LX} become typical of LLAGNs.  Additional lines of evidence which support the presence of an AGN include the fact that all three of these galaxies (NGC 2787, NGC 5866, and NGC 6500) have VLBI-detected, sub-pc scale, nuclear radio core sources (Falcke et al. 2000), a broad H$\\alpha$ component (in NGC 2787, and an ambiguous detection in NGC 5866; Ho et al.  1997b), a variable radio core in NGC 2787, and a jet-like linear structure in a high-resolution radio map at 5 GHz with the VLBA (NGC 6500; Falcke et al. 2000). Only an upper limit to the X-ray flux is obtained for NGC 5866. If an X-ray nucleus is present in this galaxy and its luminosity is only slightly below the upper limit, this source could be an AGN obscured by a column density {\\NH}$\\sim10^{23}$ {\\pcm} or larger.  If the intrinsic luminosity of the nucleus is {\\it much} lower than the observed upper limit, an AGN would have to be almost completely obscured and/or the optical emission lines powered by some other source(s). The optical classification (transition object) suggests the presence of an ionizing source other than an AGN.   \\begin{figure} \\centerline{\\psfig{figure=yuichi_terashima_fig1a.ps,width=6.8cm,angle=-90}} \\centerline{\\psfig{figure=yuichi_terashima_fig1b.ps,width=6.8cm,angle=-90}} \\caption[]{Examples of Chandra spectra. (a) NGC 3169 and (b) NGC 4548} \\end{figure}   \\subsection{Obscured LLAGNs}  In our sample, we found at least two highly absorbed LLAGNs (NGC 3169 and NGC 4548). In addition, if the X-ray faint objects discussed in the previous subsection are indeed AGNs, they are most probably highly absorbed with {\\NH}$>10^{23}$ {\\pcm}. Among these absorbed objects, NGC 2787 is classified as a LINER 1.9, NGC 3169, NGC 4548, and NGC 6500 as LINER 2s, and NGC 5866 as a transition 2 object. Thus, heavily absorbed LINER 1.9s/2s, of which few are known, are found in the present observations demonstrating that radio selection is a valuable technique for finding obscured AGNs. Along with heavily obscured LLAGNs known in low-luminosity Seyfert 2s (e.g., NGC 2273, NGC 2655, NGC 3079, NGC 4941, and NGC 5194; Terashima et al. 2002a), our observations show that at least some type 2 LLAGNs are simply low-luminosity counterparts of luminous Seyferts in which heavy absorption is often observed. Some LINER 2s (e.g., NGC 4594, Terashima et al. 2002a; NGC 4374, Finoguenov \\& Jones 2001; NGC 4486, Wilson \\& Yang 2002) and low-luminosity Seyfert 2s (NGC 3147) show no strong absorption.  Therefore, the orientation dependent unification scheme does not always apply to AGNs in the low-luminosity regime.    \\subsection{Radio Loudness of LLAGNs}   Earlier studies have suggested that LLAGNs tend to be radio loud compared to more luminous AGNs based on the spectral energy distributions of seven LLAGNs (Ho 1999) and, for a larger sample, on the conventional definition of radio loudness $R_{\\rm O}=L_{\\nu}$(5 GHz)/$L_{\\nu}$(B) (the subscript ``O'' stands for optical), with $R_{\\rm O}>10$ being radio loud (Ho \\& Peng 2001).  Ho \\& Peng (2001) measured the luminosities of the nuclei by spatial analysis of optical images obtained with {\\HST} to reduce the contribution from stellar light. A caveat in the use of optical measurements for the definition of radio loudness is extinction, which will lead to an overestimate of $R_{\\rm O}$. Although Ho \\& Peng (2001) used only type 1--1.9 objects, some objects of these types show high absorption columns in their X-ray spectra. In this subsection, we study radio loudness by comparing radio and hard X-ray luminosities. Since the unabsorbed luminosity for objects with {\\NH} \\simgt $10^{23}$ {\\pcm} (equivalent to $A_{\\rm V}$ \\simgt 50 mag for a normal gas to dust ratio) can be reliably measured in the 2--10 keV band, it is clear that replacement of optical by hard X-ray luminosity potentially yields considerable advantages.  In the following analysis, radio data at 5 GHz taken from the literature are used since fluxes at this frequency are widely available for various classes of objects.  We used the radio luminosities primarily obtained with the VLA at \\simlt $1^{\\prime\\prime}$ resolution for the present sample. High resolution VLA data at 5 GHz are not available for several objects.  For four objects among such cases, VLBA observations at 5 GHz with 150 mas resolution are published in the literature (Falcke et al. 2000) and are used here. For two objects, we estimated 5 GHz fluxes from 15 GHz data by assuming a spectral slope of $\\alpha=0$ (cf. Nagar et al. 2001). Since our sample is selected based on the presence of a compact radio core, the sample could be biased to more radio loud objects. Therefore, we constructed a larger sample by adding objects taken from the literature for which 5 GHz radio, 2--10 keV X-ray, and $R_{\\rm O}$ measurements are available.   First, we introduce the ratio $R_{\\rm X} = \\nu L_{\\nu}$(5 GHz)/{\\LX} as a measure of radio loudness and compare the ratio with the conventional $R_{\\rm O}$ parameter. The X-ray luminosity {\\LX} in the 2--10 keV band (source rest frame), corrected for absorption, is used. We examine the behavior of $R_{\\rm X}$ using samples of AGN over a wide range of luminosity, including LLAGN, the Seyfert sample of Ho \\& Peng (2001) and PG quasars which are also used in their analysis. The X-ray luminosities (mostly measured with {\\asca}) are compiled from the literature.   Fig. 2 compares the parameters $R_{\\rm O}$ and $R_{\\rm X}$ for the Seyferts and PG sample. These two parameters correlate well for most Seyferts.  Some Seyferts have higher $R_{\\rm O}$ values than indicated by most Seyferts.  This could be a result of extinction. Seyferts showing X-ray spectra absorbed by a column greater than $10^{22}$ cm$^{-2}$ (NGC 2639, 4151, 4258, 4388, 4395, 5252, and 5674) are shown as open circles in Fig. 2. At least four of them have larger $R_{\\rm O}$ than indicated by the correlation. The correlation between $\\log R_{\\rm O}$ and $\\log R_{\\rm X}$ for the less absorbed Seyferts can be described as $\\log R_{\\rm O}$ = 0.88 $\\log R_{\\rm X}$ + 5.0. According to this relation, the boundary between radio loud and radio quiet object ($\\log R_{\\rm O}$ = 1) corresponds to $\\log R_{\\rm X} = -4.5$.   The PG quasars show systematically lower $R_{\\rm O}$ values than those of Seyferts at a given $\\log R_{\\rm X}$.  For the former objects, $\\log R_{\\rm O}=1$ corresponds to $\\log R_{\\rm X}=-3.5$.  This apparently reflects a luminosity dependence of the shape of the SED: luminous objects have steeper optical-X-ray slopes $\\alpha _{\\rm ox} = 1.4-1.7$ ($S\\propto\\nu^{-\\alpha}$), where $\\alpha _{\\rm ox}$ is often measured as the spectral index between 2200 A and 2 keV, while less luminous AGNs have $\\alpha _{\\rm ox} = 1.0-1.2$ (Ho 1999).  This is related to the fact that luminous objects show a more prominent ``big blue bump'' in their spectra.  Figure 8 of Ho (1999) demonstrates that low-luminosity objects are typically 1--1.5 orders of magnitude fainter in the optical band than luminous quasars for an given X-ray luminosity.   The definition of radio loudness using the hard X-ray flux ($R_{\\rm X}$) appears to be more robust because X-rays are less affected by both extinction at optical wavelengths and the detailed shape of the blue bump. Further, measurements of nuclear X-ray fluxes are much easier than measurements of nuclear optical fluxes, since in the latter case the nuclear light must be separated from the surrounding starlight.   Fig. 3 shows the X-ray luminosity dependence of $R_{\\rm X}$. In this plot, the LLAGN sample discussed in the present paper is shown in addition to the Seyfert and PG samples used above. This is an ``X-ray version'' of the $\\log R_{\\rm O}$-$M_{B}^{\\rm nuc}$ plot (Fig. 4 in Ho \\& Peng 2001). Our plot shows that a large fraction of LLAGNs ({\\LX}$<10^{42}$ {\\eps}) are radio loud. This is a confirmation of Ho \\& Peng's (2001) finding. Since radio emission in LLAGNs is likely to be dominated by emission from jets (Nagar et al. 2001; Ulvestad \\& Ho 2001), these results suggest that, in LLAGN, the fraction of the accretion energy that powers a jet, as opposed to electromagnetic radiation, is larger than in more luminous Seyfert galaxies and quasars.     \\begin{figure} \\centerline{\\psfig{figure=yuichi_terashima_fig2.ps,width=7cm,angle=-90}} \\caption[]{ Comparison between $R_{\\rm O}=L_{\\nu}$(5 GHz)/$L_{\\nu}$(B) and $R_{\\rm X}=\\nu L_{\\nu}$(5 GHz)/{\\LX} for Seyferts and PG quasars. The conventional boundary between ``radio loud'' and ``radio quiet'' objects ($\\log R_{\\rm O}=1$) is shown as a horizontal dashed line. } \\end{figure}  \\begin{figure} \\centerline{\\psfig{figure=yuichi_terashima_fig3.ps,width=7cm,angle=-90}} \\caption[]{ X-ray luminosity dependence of $R_{\\rm X}=\\nu L_{\\nu}$(5 GHz)/{\\LX} for the present LLAGN sample, Seyfert galaxies, and PG quasars. The boundary between ``radio loud'' and ``radio quiet'' objects ($\\log R_{\\rm x}=-4.5$) is shown as a horizontal dashed line. } \\end{figure}"
    },
    "0206/astro-ph0206062_arXiv.txt": {
        "abstract": "{  We show that the first 10 eigencomponents of the Karhunen-Lo\\`eve expansion or Principal Component Analysis (PCA) provide a robust classification scheme for the identification of stars, galaxies and quasi-stellar objects from multi-band photometry. To quantify the efficiency of the method, realistic simulations are performed which match the planned Large Zenith Telescope survey. This survey is expected to provide spectral energy distributions with a resolution $R\\simeq40$ for $\\sim10^6$ galaxies to $R\\le23$ ($z\\sim 1$), $\\sim 10^4$ QSOs, and $\\sim 10^5$ stars.  We calculate that for a median signal-to-noise ratio of 6, 98\\% of stars, 100\\% of galaxies and 93\\% of QSOs are correctly classified. These values increase to 100\\% of stars, 100\\% of galaxies and 100\\% of QSOs at a median signal-to-noise ratio of 10. The 10-component PCA also allows measurement of redshifts with an accuracy of $\\sigma_\\mathrm{Res.}\\la0.05$ for galaxies with $z\\la0.7$, and to $\\sigma_\\mathrm{Res.}\\la0.2$ for QSOs with $z\\ga2$, at a median signal-to-noise ratio of 6. At a median signal-to-noise ratio 20, $\\sigma_\\mathrm{Res.}\\la0.02$ for galaxies with $z\\la1$ and for QSOs with $z\\ga2.5$ (note that for a median $S/N$ ratio of 20, the bluest/reddest objects will have a signal-to-noise ratio of $\\la 2$ in their reddest/bluest filters). This redshift accuracy is inherent to the $R\\simeq40$ resolution provided by the set of medium-band filters used by the Large Zenith Telescope survey. It provides an accuracy improvement of nearly an order of magnitude over the photometric redshifts obtained from broad-band $BVRI$ photometry. ",
        "introduction": "\\label{intro}  The galaxy luminosity function (hereafter LF), defined as the number density of galaxies per unit interval of luminosity is a fundamental statistical tool required to model the formation, evolution and clustering of the galaxies. At $z\\la1$, it is well established that the LF depends on the galaxy morphological type \\citep{bingelli88,marzke98,loveday99}, and that it evolves with redshift \\citep{lilly96,lin99,bromley98,lapparent02a}. Measurement of the LF thus requires large galaxy samples which can be separated into several morphological, spectral or color classes and in redshift intervals. Among the next generation redshift surveys, only 3 will be able to probe the galaxy LF to $z\\sim1$. The DEEP redshift survey using Keck telescopes \\citep{davis98} and the VIRMOS redshift survey using the VLT \\citep{lefevre98} are optimized for clustering analyses, but they will also provide measurements of the galaxy LFs, even if the 2 surveys shall suffer from various aperture effects and calibration difficulties. The third survey is the Large Zenith Telescope (hereafter LZT) survey, which is optimized for the measurements of the galaxy LFs to $z\\sim1$.  An essential step in the measurement of the LF for a systematic survey is the classification of objects as stars, galaxies, QSOs, etc. For the galaxies and QSOs, an estimate of the redshift for each object is also desired. This paper proposes and tests a classification approach based on the Principal Component Analysis (PCA), also known as Karhunen-Lo\\`eve expansion -- the underlying principles of the PCA were independently derived by \\citet{karhunen47} and \\citet{loeve48}. The PCA is a non-parametric approach which has been successfully used for a variety of astronomical applications including stellar classification from photometric data \\citep{deeming64,scarfe66,whitney83,whitney83b} and from spectra \\citep{storri94,ibata97,bailer98,singh98}, galaxy classification from photometric data \\citep{watanabe85} and from galaxy spectra \\citep{connolly95a,connolly95b,galaz98,connolly99,ronen99}, and for galaxy redshift measurements \\citep{glazebrook98}; other fields of application are solar flare observations \\citep{teuber79}, asteroid spectra \\citep{britt92}, inter-stellar medium emission lines \\citep{heyer97}, gamma ray bursts \\citep{bagoly98}, and active galaxies  \\citep{mittaz90,dultzin96,turler98}.  All previous spectral classification attempts using the PCA employed either multicolor photometry (usually fewer than 10 color bins, e.g. $UBVRIJHK$) or medium to high-resolution spectroscopy (resolution $R > 500$). The PCA has not been tested on spectral energy distributions (SED) with $R \\sim 40$  because no such data existed until the UBC-NASA Multi-narrowband survey \\citep{hickson98a,cabanac98}.  In this paper we use simulations based on the LZT survey parameters to evaluate the PCA method. Section \\ref{simul} describes our simulations of mock LZT catalogues, Sect{.} \\ref{method} describes the approach used to classify the objects using the PCA, Sect{.} \\ref{result} shows the efficiency of the classification and Sect{.} \\ref{redshift} discusses the redshift accuracies which can be obtained directly from the PCA or from a composite method similar to that described by \\citet{glazebrook98}. ",
        "conclusions": "This paper describes an application of Principal Component Analysis (PCA) to a simulated multicolor survey using the 40 medium-band filters of the Large Zenith Telescope. For that purpose, we generate realistic mock catalogues of $\\sim3000$ stars, $\\sim 30\\,000$ galaxies, and $\\sim 1\\,000$ QSOs. For stars, we use templates from the library of Pickles (1998) and the phenomenological model of star counts of Bahcall \\& Soneira (1986). For galaxies, we use spectral energy distributions (SED) from GISSEL (Charlot, 1993) and PEGASE (Fioc, 1997) and a luminosity function derived from a review of the most recent $R$-band luminosity functions of the literature. We choose the CFRS-type evolution in the galaxy luminosity function, with luminosity evolution of only the late-type galaxies. For QSOs, we use an extrapolation of the composite spectrum of Francis (1991), and the luminosity function of the 2dF QSO survey.  Using the realistic mock catalogues, we perform a PCA and extract the first 10 eigencomponents. The 10-D space allows one to separate efficiently stars, galaxies, and QSOs even at low signal-to-noise ratios.  98\\% of stars, 100\\% of galaxies and 93\\% of QSOs are classified correctly at a median signal-to-noise ratio $S/N=6$. These values increase to 100\\% of stars, 100\\% of galaxies and 100\\% of QSOs at a median $S/N=10$. For SEDs with a median $S/N=6$, the 10-component PCA also provides a measurement of redshifts accurate to $\\sigma_\\mathrm{Res.}\\la0.05$ for galaxies with $z\\la0.7$, and to $\\sigma_\\mathrm{Res.}\\la0.2$ for QSOs with $z\\ga2$. At a median $S/N=20$, $\\sigma_\\mathrm{Res.}\\la0.02$ for galaxies with $z\\la1$ and for QSOs with $z\\ga2.5$ (for a given median $S/N$, a 10 to 30 times lower $S/N$ is expected at the extreme wavelengths of the bluest/reddest objects). This is not sufficient for small-scale 3-D clustering analyses, but perfectly adequate for luminosity function studies, and for measuring the evolution with redshift in the large-scale clustering using projected moments.  This paper also underlines the main weakness of the PCA. It is well-known that age, star-formation rate, redshift, and dust extinction produce degenerate SEDs at a resolution $R\\simeq40$. Although the PCA is not able to resolve some intrinsic degeneracies due to the medium-band observing technique, it efficiently reduces the noise in the SEDs at the expense of additional degeneracies. The solution to this problem may lie in the combination of a PCA with a standard $\\chi^2$ fitting procedure. Another crucial issue in the use of the PCA for type/class/redshift measurement, is to calculate the eigenvectors from a sample in which each type of object is sufficiently well represented. The use of such a catalog, constructed from a combination of a wide variety of well calibrated observed SEDs together with precise evolutionary models, will guarantee the best results for PCA analyses. These reference samples will also allow detection of new types of object, as these will significantly deviate from the sequences of known objects."
    },
    "0206/astro-ph0206254_arXiv.txt": {
        "abstract": "We describe \\boomn; a balloon-borne microwave telescope designed to map the Cosmic Microwave Background (CMB) at a resolution of 10$'$ from the Long Duration Balloon (LDB) platform.   The millimeter-wave receiver employs new technology in bolometers, readout electronics, cold re-imaging optics, millimeter-wave filters, and cryogenics to obtain high sensitivity to CMB anisotropy.  Sixteen detectors observe in 4 spectral bands centered at 90, 150, 240 and 410 GHz. The wide frequency coverage, the long duration flight, the optical design and the observing strategy provide strong rejection of systematic effects.  We report the flight performance of  the instrument during a 10.5 day stratospheric balloon flight launched from McMurdo Station, Antarctica that mapped $\\sim$ 2000 square  degrees of the sky. ",
        "introduction": "Measurements of the angular power spectrum of anisotropy in the Cosmic Microwave Background (CMB) are greatly enhancing our knowledge of fundamental properties of the universe.   In particular, models of the early universe predict the existence of a series of harmonic peaks in the angular power spectrum at degree scales.  The precise determination of the amplitude and angular scale of the peaks provides strong constraints on these models and enables their parameters to be determined with great precision (\\cite{sachswolfe}; \\cite{hu}).  After the discovery of the large scale anisotropies by COBE-DMR (\\cite{Bennett}), a host of ground-based and balloon-borne  observations detected a first peak at a multipole moment of $\\ell \\sim$ 200 in the angular power spectrum of the CMB (\\cite{bjk};\\cite{MAT};\\cite{b97}). Here we describe the \\boom experiment, which has made a deep map of the CMB at sub-degree resolution and a precision measurement of the angular scale and amplitude of the first peak (\\cite{pdb}; \\cite{pdb02}) from the Long  Duration Balloon (LDB) platform.  The \\boom instrument consists of a 1.3-m off-axis telescope that feeds a bolometric array receiver.  The receiver is housed inside a long duration liquid helium cryostat. A sorption pumped $^3$He refrigerator maintains the detectors at 280 mK. Observations are made in 4 spectral bands centered at 90, 150, 240, and 410 GHz (3mm, 2mm, 1.3mm, and 750 $\\mu$m) with angular resolutions of 18$'$, 10$'$, 14$'$, and 13$'$ FWHM respectively. A test flight of the \\boom payload in a different configuration (\\cite{b97inst}) flew in a 6 hour engineering flight from Palestine, Texas in 1997 (\\cite{b97}; \\cite{b972}).  The configuration of the instrument described here flew in a 252 hour flight from McMurdo Station, Antarctica in 1998-1999. ",
        "conclusions": "The \\boom payload performed to specification in its first long duration flight above Antarctica.  The cryogenic system held the detectors below 0.285~K for the duration of the 10.5 day flight. The attitude control system and the data acquisition performed flawlessly. Ambient temperatures on the payload stayed within operating range.  The receiver performed well, achieving an average instantaneous sensitivity of 140 $\\mu$K$\\sqrt{\\rm s}$ in the 90 and 150~GHz single mode channels.  The micromesh design for the bolometer absorbers limited the cosmic ray contamination to 5\\% of the data.  The quasi total power radiometer provided stability over a wide enough range of signal frequencies to allow mapping of roughly 2000 square degrees of sky, a calibration to 10\\% from the CMB dipole, and precision measurement of CMB anisotropies on degree scales.  The cosmological results from the Antarctic flight of this instrument are reported in \\cite{pdb}, \\cite{b98parm}, \\cite{pdb02}, and \\cite{cbn2}.  \\smallskip  The \\boom project is supported by the CIAR and NSERC in Canada; by PNRA, Universit\\'{a} ``La Sapienza'' and ASI in Italy; by PPARC in the UK; and by NASA, NSF, OPP, and NERSC in the US.  The authors would like to thank Kathy Deniston for logistical support, and NASA's National Scientific Balloon Facility (NSBF) and the US Antarctic Program for excellent field and flight support."
    },
    "0206/astro-ph0206312_arXiv.txt": {
        "abstract": "{ We have detected diffuse soft X-ray emission (0.4$-$1 keV) from the disk of the spiral galaxy NGC~2403 with {\\it Chandra}. This diffuse emission (with a total luminosity of 2.1 $\\times$ 10$^{38}$ erg s$^{-1}$ and a gas temperature of 2$-$8 $\\times$ 10$^{6}$ K) is well separated from the numerous bright point sources. NGC~2403 is a luminous spiral galaxy with a high rate of star formation. } Recent H~{\\small I} observations have revealed an extended H~{\\small I} halo with anomalous velocities and a general inflow towards the central regions of the galaxy. This result and the present detection of a diffuse, hot X-ray emitting gas point at a very active disk-halo connection and galactic fountain types of phenomena. ",
        "introduction": "Before the launch of {\\it Chandra} the study of diffuse X-ray emission from normal (non-starburst) spiral galaxies was seriously impeded by lack of spatial resolution \\citep{fab89}. The point sources could not be well separated from the diffuse thermal gas. The detection of diffuse emission from the disk of spiral galaxies was therefore only possible for some nearby objects (e.g.\\ M~33: \\citet{lon96}; M~101: \\citet{sno95}; M~51: \\citet{ehl95}). Coronal (halo) emission was observed in some nearby edge-on spiral galaxies like NGC~891 \\citep{bre97} and NGC~4631 \\citep{wan95, wan01}. Also the diffuse component of the hot thermal plasma in the Milky Way has been studied extensively (e.g. \\citet{kan97, val98}).  Here we present {\\it Chandra} observations of the spiral galaxy NGC~2403 that made it possible to unambiguously separate discrete sources from diffuse emission. This galaxy is a nearby Sc spiral, morphologically similar to M~33, viewed at an inclination angle of $\\sim$~60$^{\\circ}$. It is well isolated on the sky and shows no signs of recent interactions. H~{\\small I} observations \\citep{fra02} have revealed a kinematically anomalous component of neutral gas. The H~{\\small I} position-velocity (p-v) diagram along the major axis of this galaxy (Figure~\\ref{fig1}) shows systematic asymmetries in the form of wings in the line profiles at lower velocities with respect to the rotation curve (white squares). In the central part of the galaxy such wings extend up to 150 km s$^{-1}$ from the rotation curve. This kinematical pattern is different from that expected for a thin cold disk of H~{\\small I} (see the model in the right panel of Figure~\\ref{fig1}). This anomalous gas was previously unknown. It has not been detected before, mainly because of lack of sensitivity of the observations.  A detailed analysis of the H~{\\small I} data of NGC~2403 has shown that the wings of H~{\\small I} are produced by gas (the ``anomalous gas'') located above the plane of the galaxy and rotating $\\sim$~20$-$50 km~s$^{-1}$ more slowly than the gas in the disk. A similar slower rotation of the ``halo'' gas had already been observed in the edge-on galaxy NGC~891 \\citep{swa97}. The anomalous gas in NGC~2403 extends out to $\\sim$ 15 kpc from the centre of the galaxy and has a total mass of about 3 $\\times$ 10$^8$ $M_{\\odot}$ ($\\sim$~1/10 of the total H~{\\small I} mass). The study of its velocity field has also revealed a probable large-scale radial inflow (10$-$20 km~s$^{-1}$) towards the centre of the galaxy \\citep{fra01}.  A possible interpretation is that of a {\\it galactic fountain} \\citep{sha76, bre80} and the observed neutral gas may be the result of cooling of ionised gas blown up from the disk into the halo. The discovery of hot X-ray emitting gas from the disk of NGC~2403, reported here, supports this picture. The anomalous H~{\\small I} may be related to the hot X-ray gas. ",
        "conclusions": "We have detected diffuse soft X-ray emission from the disk of the spiral galaxy NGC~2403. The X-ray emission is strongly associated with the H~{\\small II} regions and is most likely produced by hot gas from stellar winds and expanding supernova shells within the disk of the galaxy. If we assume a rate of supernova explosions of 0.01 events yr$^{-1}$ for NGC~2403 \\citep{mat97} and a mean energy release per supernova of 10$^{51}$ erg we estimate that the total amount of energy released by supernovae is $\\sim$~3 $\\times$ 10$^{41}$ erg s$^{-1}$. This is much larger than the observed X-ray luminosity (2.1 $\\times$ 10$^{38}$ erg s$^{-1}$). Thus the hot ISM of NGC~2403 can be easily produced by energy input from star formation activity and a heating efficiency as low as 0.001.  One interesting question is whether the hot gas in NGC~2403 is actually bound to the galaxy or not. The escape velocity from a galaxy is usually calculated as $v_e(r)=\\sqrt{2|\\phi(r)|}$ where $\\phi(r)$ is the potential at radius $r$. Assuming a mass model (disk $+$ isothermal dark matter halo) for NGC~2403 \\citep{fra02} we get an escape velocity of 300$-$350 km s$^{-1}$ in the central 4 kpc. The thermal velocity of the hot gas is likely to be between the adiabatic sound speed $v_{so}=\\sqrt{\\frac{5 kT}{3 \\mu m_H}}$ and the post-shock speed $v_{sh}=\\sqrt{\\frac{16 kT}{3 \\mu m_H}}$. If we take the temperature of the one-component Mekal fit kT $\\simeq$ 0.25 keV (T $\\simeq$ 2.9 $\\times$ 10$^6$ K) we obtain a thermal velocity for the hot gas (n$_{He}$/n$_{H}$=0.1) of 170$-$300  km s$^{-1}$, lower than the escape velocity. Similar values have been found for other galaxies (e.g. NGC~4631, \\citet{wan95} and the Milky Way, \\citet{kan97}). The above estimate shows that at least part of the X-ray emitting gas found in NGC~2403 is likely to be bound to the galaxy. One implication of this result is that this hot gas does not substantially contribute to the enrichment of the surrounding intergalactic medium. Furthermore, if this gas can temporarily escape from the disk into the halo region it will necessarily fall down onto the disk. According to a ``galactic fountain'' type of process the hot gas escaping from the disk is expected to cool quickly because of adiabatic expansion and in the final phase it could become observable as neutral gas at anomalous velocities. The anomalous H~{\\small I} shown in the p-v diagram of Figure~\\ref{fig1} is likely to represent such a final phase of the galactic fountain, whereas the hot X-ray emitting gas would pertain to the initial phase of it.  From the above X-ray luminosity and temperature we estimate an electron density and a mass of the hot gas of 0.15 $\\times$ 10$^{-3}$ $h^{-\\frac{1}{2}} f^{-\\frac{1}{2}}$ cm$^{-3}$ and 6.0 $\\times$ 10$^{6} h^{\\frac{1}{2}} f^{-\\frac{1}{2}}$ $M_{\\odot}$ where $h$ is the scaleheight of the disk in kpc and $f$ is the filling factor. The cooling rate of the hot gas can be estimated with the formula \\citep{nul84} $\\dot{M}_{cool} \\approx \\frac{M_{X}}{t_{cool}} = \\frac{2 \\mu m_H L_X}{3 k T}$ that for $n_{He}/n_{H}=0.1$ gives $\\dot{M}_{cool} \\sim 0.01$ $M_{\\odot}$ yr$^{-1}$. If we consider now the mass of the anomalous H~{\\small I}, a typical size of 10 kpc and the measured infall velocity of 10$-$20 km s$^{-1}$, we estimate the infalling rate of H~{\\small I} to be $\\dot{M}_{HI} \\sim 0.3-0.6$ $M_{\\odot}$ yr$^{-1}$ larger than the cooling rate of the hot gas. However such value for cooling rate is probably representative only for the disk component and would vary substantially once the gas has left the disk itself. A more useful quantity is the outflowing rate of the hot gas. Considering the mass and the thermal velocity determined above, and assuming a continous cycle beetwen hot and neutral gas, such escaping rate turns out to be $\\dot{M}_{hot} \\sim$ 0.1$-$0.2 $M_{\\odot}$ yr$^{-1}$, of the same order as the infalling H~{\\small I}. Therefore it is possible that cooling of the X-ray emitting gas in NGC~2403 and its motion via a galactic fountain could produce the observed H~{\\small I} position-velocity pattern.  Is has been suggested that the anomalous H~{\\small I} is common among spiral galaxies \\citep{fra01}. Similarly, the hot X-ray emitting gas is probably present in several spiral galaxies with a high rate of star formation. One can argue, therefore, that the hot gas and the anomalous H~{\\small I} are, generally, connected as in NGC~2403. Finally this discovery of diffuse X-ray emitting gas in NGC~2403 has also interesting implications for the study of our Galaxy. The anomalous H~{\\small I} in NGC~2403 is analogous to at least some of the High Velocity Clouds \\citep{wak97} of our Galaxy. The suggested relation between the neutral and the hot gas in NGC~2403 supports the possibility that galactic fountains play an important role also in the Milky Way."
    },
    "0206/astro-ph0206318_arXiv.txt": {
        "abstract": "We complete construction of a catalog containing improved astrometry and new optical/infrared photometry for the vast majority of NLTT stars lying in the overlap of regions covered by POSS I and by the second incremental 2MASS release, approximately 44\\% of the sky.  The epoch 2000 positions are typically accurate to 130 mas, the proper motions to $5.5\\,\\masyr$, and the $V-J$ colors to 0.25 mag.  Relative proper motions of binary components are meaured to $3\\,\\masyr$.  The false identification rate is $\\sim 1\\%$ for $11\\la V\\la 18$ and substantially less at brighter magnitudes. These improvements permit the construction of a reduced proper motion diagram that, for the first time, allows one to classify NLTT stars into main-sequence (MS) stars, subdwarfs (SDs), and white dwarfs (WDs). We in turn use this diagram to analyze the properties of both our catalog and the NLTT catalog on which it is based.  In sharp contrast to popular belief, we find that NLTT incompleteness in the plane is almost completely concentrated in MS stars, and that SDs and WDs are detected almost uniformly over the sky $\\delta>-33^\\circ$.  Our catalog will therefore provide a powerful tool to probe these populations statistically, as well as to reliably identify individual SDs and WDs. ",
        "introduction": "}  In Paper I \\citep{paper1}, we discussed the motivation for improving the astrometry and the photometry of the $\\sim 59,000$ high proper-motion ($\\mu\\geq 180\\,\\masyr$) stars in the {\\it New Luyten Two Tenths} (NLTT) Catalog \\citep{luy}, and we outlined our basic strategy for obtaining these improvements: at the bright end, match NLTT star to entries in the Hipparcos \\citep{hip}, Tycho-2 \\citep{t2}, and Starnet \\citep{starnet} catalogs and, at the faint end, find counterparts of USNO-A \\citep{usnoa1,usnoa2} stars in 2MASS \\citep{2mass} whose position offsets are predicted by the proper motions listed in NLTT.  We carried out the bright-end search and used the results to characterize the position and proper-motion (PPM) errors of the NLTT.  Here we complete the catalog (in the regions covered by the first 2MASS release and the first Palomar Observatory Sky Survey -- POSS I) by carrying out the faint-end search, and we integrate the results of both searches. This paper therefore has three interrelated goals. First, to give a comprehensive account of how the catalog was constructed and of the various tests we performed to determine the accuracy of our identifications and our proper-motion measurements. Second, to characterize the properties of the original NLTT catalog, now that much better astrometry and photometry are available for a large fraction of its stars.  Third, to characterize the properties of the catalog we have constructed, including its precision and its completeness relative to NLTT.  We expect that the importance of the first and third goals are obvious to the reader, but the reasons for characterizing the NLTT just at the time when it is being superseded may not be.  First, in constructing the present catalog, we made heavy use of all aspects of NLTT including not only its positions, proper motions, and photometry, but also its extensive notes on binaries.  Our approach to carrying out this work was heavily influenced by our assessment of NLTT's properties, and hence ours catalog's construction and limitations cannot be fully understood without a knowledge of these properties.  Second, the publication of NLTT was a watershed in the history of astronomy: the lengthy list of papers that we cited in Paper I comprise but a tiny fraction of the literature that is based directly or indirectly on NLTT.  A good understanding of the NLTT's properties will be useful in assessing which of the conclusions drawn by these papers can be relied upon, and which require further refinement.  Third, more work revising the NLTT remains to be done.  For technical reasons that will be described below, our catalog will cover only the portion of the sky $\\delta > -33^\\circ$. And, of course, as of now it is restricted to the $\\sim 1/2$ of sky covered by the 2MASS second incremental release. Hence, for the present, the original NLTT remains the best source of information for 23,000 of its 59,000 stars. ",
        "conclusions": "\\label{conclude}}  The catalog presented here gives improved astrometry and photometry for the great majority of stars in the NLTT that lie in the overlap of the areas covered by the second incremental 2MASS release and those covered by POSS I (basically $\\delta>-33^\\circ$).  In addition, essentially all bright NLTT stars over the whole sky have been located in PPM catalogs and, whenever possible, the close binaries among them have been resolved using TDSC.  We recover essentially 100\\% of NLTT stars $V<10$, about 97\\% for $10<V<18$, and a declining fraction thereafter.  Most of the faint stars that we do not recover are WDs, which are too faint and too blue to show up in one or both of USNO-A or 2MASS, the two catalogs that we combine to obtain our astrometry and photometry.  The catalog contains a total of 36,020 entries for 35,662 NLTT stars including a total of 723 entries for 361 NLTT ``stars'' that have been resolved in TDSC.  Of these 36,020 entries, 1353 ($\\sim 4\\%$) contain new positions and photometry but not new proper motions because they have been identified in only USNO-A or 2MASS, but not both.  The new positions are accurate to 130 mas. The new proper motions are accurate to $5.5\\,\\masyr$, more than a 3-fold improvement over NLTT. Narrow angle proper motions are accurate to $3\\,\\masyr$.  The catalog provides a powerful powerful means to investigate SDs, WDs, faint MS stars as well as CPM binaries."
    },
    "0206/astro-ph0206404_arXiv.txt": {
        "abstract": "Recent X-ray observations of galaxies with ROSAT, ASCA, and Chandra have revealed numerous bright off-center point sources which, if isotropic emitters, are likely to be intermediate-mass black holes, with $M\\sim 10^{2-4}\\,M_\\odot$.  The origin of these objects is under debate, but observations suggest that a significant number of them currently reside in young high-density stellar clusters.  There is also growing evidence that some Galactic globular clusters harbor black holes of similar mass, from observations of stellar kinematics. In such high-density stellar environments, the interactions of intermediate-mass black holes are promising sources of gravitational waves for ground-based and space-based detectors. Here we explore the signal strengths of binaries containing intermediate-mass black holes or stellar-mass black holes in dense stellar clusters. We estimate that a few to tens per year of these objects will be detectable during the last phase of their inspiral with the advanced LIGO detector, and up to tens per year will be seen during merger, depending on the spins of the black holes.  We also find that if these objects reside in globular clusters then tens of sources will be detectable with LISA from the Galactic globular system in a five year integration, and similar numbers will be detectable from more distant galaxies.  The signal strength depends on the eccentricity distribution, but we show that there is promise for strong detection of pericenter precession and Lense-Thirring precession of the orbital plane.  We conclude by discussing what could be learned about binaries, dense stellar systems, and strong gravity if such signals are detected. ",
        "introduction": "We are entering an era in which numerous experiments will search for astrophysical gravitational radiation.  Ground-based detectors such as LIGO (Barish 2000), VIRGO (Fidecaro et al. 1997), GEO600  (Schilling 1998), TAMA (Ando et al. 2002), and others focus on relatively high frequencies $f_{\\rm GW}\\approx 10-1000$~Hz.  These frequencies are appropriate for the final inspiral and merger of binaries with total masses $\\sim 1-1000\\,M_\\odot$. Space-based missions such as the Laser Interferometer Space Antenna (LISA; see, e.g., Danzmann 2000) will complement this frequency range by focusing on comparatively low frequencies, $f_{\\rm GW}=10^{-4}-1$~Hz. In this range there are a number of ``guaranteed\" sources in the form of main sequence and white dwarf binaries, but there is much  more uncertainty about the prevalence and properties of double compact binaries containing neutron stars and black holes.  Such binaries are of great interest for many reasons, such as their use as probes of strong gravity and their potential to illuminate other aspects of astronomy, including the evolution of galaxies.  The existence of a new class of double compact binaries is suggested by two  recent lines of evidence.  First, observations of the kinematics of the central regions of some globular clusters suggests that black holes with $M\\gta 100\\,M_\\odot$ may exist in their cores (e.g., Gebhardt et al. 2000; D'Amico et al. 2002).  Second, X-ray observations have revealed variable, unusually high flux point sources in a number of galaxies (e.g., Fabbiano 1989; Fabbiano, Schweizer, \\& Mackie 1997; Colbert \\& Mushotzky 1999; Zezas, Georgantopoulos, \\& Ward 1999; Kaaret et al. 2001; Matsumoto et al. 2001; Fabbiano, Zezas, \\& Murray 2001; see Colbert \\& Ptak 2002 for a catalog of objects).  The rapid and strong variability of these sources indicates that they are black holes.  If their flux is not strongly beamed, then for their luminosities to be below the Eddington luminosity $L_E=1.3\\times 10^{38}(M/M_\\odot)$~erg~s$^{-1}$, at which radiation forces balance gravity, the masses of the brightest sources must be at least $\\sim 10^3\\,M_\\odot$ (e.g., Matsumoto et al. 2001).  In addition, their off-center positions in their host galaxies indicate that in many cases the masses cannot be substantially greater than $\\sim 10^5\\,M_\\odot$, otherwise dynamical friction would cause the black holes to sink rapidly to the dynamical center (Kaaret et al. 2001; Tremaine, Ostriker, \\& Spitzer 1975). Recent optical observations suggest that for at least some of the sources, the X-rays may not be beamed strongly, based on recombination emission from surrounding nebulae (Pakull \\& Merioni 2002). This implies that the total luminosity is indeed high, so that if the sources are sub-Eddington (but see Begelman 2002) there may therefore be a significant number of $\\sim 10^{2-4}\\,M_\\odot$ black holes in the universe. Currently, these sources are preferentially in star forming regions or in starburst galaxies (Matsumoto et al. 2001; Liu \\& Bregman 2001).  A number of these objects are associated with dense young stellar clusters (e.g., Matsumoto et al. 2001).  Combined with the kinematic evidence in globular clusters, this implies that black holes of this sort undergo frequent dynamical encounters.  It has been suggested that the ultraluminous X-ray sources originated due to dynamical encounters in globular clusters, then merged with their host galaxies (Miller \\& Hamilton 2002a,b).  Other proposals, if these objects are $10^{2-4}\\,M_\\odot$ black holes instead of being strongly beamed (King et al. 2001; King 2002; K\\\"ording, Falcke, \\& Markoff 2002) or super-Eddington by a large factor (Begelman 2002), are that they form as a result of a core collapse of a massive young star cluster (Matsushita et al. 2000; Ebisuzaki et al. 2001) or as remnants of Population~III stars (Madau \\& Rees 2001).  Here we consider the implications for gravitational wave emission if these sources are indeed intermediate-mass black holes.  Independent of  their detailed origin, black holes in dense stellar environments are promising sources of gravitational radiation due to the diversity of dynamical interactions that are possible (see also Benacquista 1999, 2002a,b; Benacquista, Portegies Zwart, \\& Rasio 2001).   We will focus mainly on black holes in globular clusters, but will also address signals that may be evident from young stellar clusters, regardless of whether the black hole formed there or elsewhere.  In \\S~2 we give general arguments about the properties of these sources, based on the ages of globulars and the number and types of secondaries.  In \\S~3 we estimate the rates of encounters, and in \\S~4 we present calculations of the signal to noise ratios expected for LISA from Galactic globulars and from globulars in the Virgo cluster, and for coalescence of black holes as observed with the LIGO-II detector.  In \\S~5 we discuss what information could be gleaned from a reasonably high signal to noise detection.  We show that both pericenter precession and Lense-Thirring precession are potentially measurable.  These effects may allow preliminary mapping of the spacetime around black holes and could even yield an independent measure of the distance to the Virgo cluster. We present our conclusions and discuss the uncertainties in our estimates in \\S~6. ",
        "conclusions": "Dense stellar clusters are promising locations for sources of gravitational radiation.  We estimate that up to several tens of times per year, the advanced LIGO detector will see the coalescence of a small black hole with a larger one (a few to tens during inspiral, and most of the rest during merger, with both numbers dependent on black hole spin).   Mergers of two stellar-mass black holes in clusters will likely be detectable at a rate of a few per year. As always in astrophysical scenarios there are sources of uncertainty. The most major is the question of the number of black holes that are to be expected in a dense cluster, whether in the $\\sim 10\\,M_\\odot$ stellar-mass range or the $\\sim 10^{2-4}\\,M_\\odot$ intermediate-mass range.  Perhaps the most easily calculable input to this question is the number of black holes that are originally produced.  If the initial mass function for the cluster was the Salpeter function $dN/dM\\propto M^{-2.35}$ above $1\\,M_\\odot$ and flatter below $1\\,M_\\odot$ (e.g., Meyer et al. 2000), then approximately 0.1\\% of stars began with $M>25\\,M_\\odot$ and presumably evolved into black holes.  If the initial mass function was instead the Scalo (1986) distribution, which drops off more sharply at higher masses, the fraction with $M>25\\,M_\\odot$ is more like 0.05\\% (Portegies Zwart \\& McMillan 2000).  Current dense clusters have $\\sim 10^6$ stars, and may have had several times more at birth, so the number of stars that evolved into black holes is typically $\\sim 10^3$.  There are, however, numerous ways in which black holes may be lost from the cluster.  The first is in the supernova that produced the black hole. Neutron stars are known to have significant birth kicks of tens to hundreds of km~s$^{-1}$ (e.g., Hansen \\& Phinney 1997; Fryer \\& Kalogera 1997).  The mechanism for this is still debated (Spruit \\& Phinney 1998; Kusenko \\& Segre 1999;  Lai, Chernoff, \\& Cordes 2001), but it is thought that similar kick velocities for black holes are much smaller (e.g., Brandt, Podsiadlowski, \\& Sigurdsson 1995; Podsiadlowski et al. 2002; see Nelemans, Tauris, \\& van den Heuvel 1999 for a somewhat different perspective), if for no other reason than that black holes are several times more massive than neutron stars so that a fixed energy or momentum in the kick would lead to reduced speeds.  It is therefore plausible that black holes do not receive birth kicks of $\\gta 50$~km~s$^{-1}$, in which case they are retained in the cluster.  A second loss mechanism involves three-body recoil.  Several simulations have shown that black holes of a fixed mass $\\sim 10\\,M_\\odot$ in binaries with other such black holes tend to be ejected by three-body interactions before they can merge by gravitational radiation (Kulkarni et al. 1993; Sigurdsson \\& Hernquist 1993; Portegies Zwart \\& McMillan 2000).  The actual fraction of ejected black holes can depend on the mass function of stars and other variables, but recent estimates suggest that 10\\% or more of the initial black holes can be retained by their clusters over a Hubble time (Portegies Zwart \\& McMillan 2000).  If more massive black holes are present initially, then Miller \\& Hamilton (2002a) showed that they are usually not ejected.  These can grow by mergers after multiple three-body encounters, but typically they will eject several to tens of field stars  along the way.  The majority of the encounters, however, will not be with black holes so this mechanism is not expected to deplete the black hole supply significantly.  There are, in addition, at least two ways in which multibody interactions can produce a merger without accompanying dynamical recoil.  One, discussed by Miller \\& Hamilton (2002b), is that binary-binary interactions can produce a stable hierarchical triple system, and if the inclination of the orbit of the tertiary to the orbit of the inner binary is in the right range then a secular Kozai resonance can increase the eccentricity of the inner binary to the extent that it merges before the next encounter, without significant recoil. The impact of this effect depends on the binary fraction and the distribution of inclinations (see Miller \\& Hamilton 2002b), but this could allow the merger of some tens of percent of the original population of black holes.  The other recoilless possibility involves resonant encounters in three-body interactions, in which the three objects orbit hundreds or thousands of times before resolving into a binary and an unbound single star.  If, during these orbits, two black holes pass close enough to each other that losses to gravitational radiation in a single pass cause rapid merger, then again there is no dynamical kick. Various estimates suggest that in an equal-mass three-body encounter the probability of the closest approach being  less than $\\epsilon<1$ times the initial semimajor axis is $\\sim \\epsilon^{1/2}$ (Hut 1984; McMillan 1986; Sigurdsson \\& Phinney 1993).  Two $10\\,M_\\odot$ black holes must approach to within $\\sim 10^9$~cm to merge in a year, which will happen in a given resonant encounter with a probability of a few tenths of a percent for a semimajor axis of a few astronomical units.  If it takes $\\sim 10$ equal-mass encounters to harden a binary to the point of ejection, this suggests that several percent of binaries will merge before ejection in this fashion. Combining all of the above effects, it seems likely that tens of percent of the original black hole population will not be ejected by three-body recoil, leaving a present-day population of hundreds.  The third loss mechanism involves the merger itself.  The emission of gravitational waves during inspiral and merger will be somewhat asymmetric, leading to recoil.  Calculations thus far have focused on the weak-field limit.  They suggest that in the post-Newtonian limit the kick scales as $a_{\\rm LSO}^{-4}$ with the separation  $a_{\\rm LSO}$ at the last orbit before dynamical instability. At the separation $a_{\\rm LSO}=6GM/c^2$ appropriate for test particles around slowly rotating black holes, the kick will be a few kilometers per second for the mass ratio $M/m=2.6$ that maximizes the kick (Fitchett 1983; Wiseman 1992). Binaries with large mass ratios or nearly equal masses experience less recoil (for example, by  symmetry, equal-mass binaries have no kick).  Also, $a_{\\rm LSO}$ is greater in a comparable-mass binary than in the test particle limit (Pfeiffer et al. 2000), which also decreases the kick.  It is not clear what level of recoil is to be expected in the merger phase, where the radial velocity becomes rapid.  Strong-field calculations are required to resolve whether this process is dominant (perhaps kicking most merging black holes out of the cluster) or insignificant (if the recoil speeds are much less than the $\\sim 50$~km~s$^{-1}$ escape speeds from the core).  Thus, if mergers do not kick black holes out of dense clusters, one may expect at least tens to hundreds of black holes in many current systems. These are expected to reside primarily in the core of the cluster, where they have a greater tendency to interact with the more massive (and hence compact) objects in the cluster.  This may explain why no definitive examples of black hole low-mass X-ray binaries are known in the globular cluster systems of the Milky Way or Andromeda; such a population would not undergo mass transfer and would thus be observable only by its gravitational wave emission. Note, however, that there is a population of $>10^{39}$~erg~s$^{-1}$ sources in the globulars around a number of ellipticals (Angelini, Loewenstein, \\& Mushotzky 2001; White, Sarazin, \\& Kulkarni 2002; Kundu, Maccarone, \\& Zepf 2002).  Possible differences between these systems are an important subject for future study.  What about the fraction of clusters with intermediate-mass black holes? Miller \\& Hamilton (2002a) estimate that clusters with central densities greater than $\\sim 10^5$~pc$^{-3}$ have high enough encounter rates to produce $10^{2-4}\\,M_\\odot$ black holes.  In the Milky Way globular system, this would imply that roughly 40\\% of globulars could host such objects (Pryor \\& Meylan 1993). However, to be conservative, we have adopted 10\\% as our fiducial value for the estimates of merger rates.  We have also been conservative in assuming that the number density and mass of globulars is the same out to 2-3~Gpc as it is in the local universe.  As first discussed by Aguilar, Hut, Ostriker (1988), evaporation and tidal interactions attenuate the globular system of a galaxy.  Therefore, it is possible that coalescence rates a few billion years ago were higher by up to a factor of a few than they are now (Portegies Zwart \\& McMillan 2000), but this is highly uncertain.  The general model described here is one that can be tested and enhanced in ways both observational and theoretical.  From the observational standpoint it is important to continue kinematic work on globulars to look for evidence of the velocity and density cusps that are expected signatures of black holes (Bahcall \\& Wolf 1975; Frank \\& Rees 1976; see Gebhardt et al. 2000, D'Amico et al. 2002 for recent results). Further characterization of the ultraluminous X-ray sources is also important. For example, if a mass estimate can be obtained via radial velocity measurements of a companion, this will shed new light on the nature of these objects.  From the theoretical standpoint there are several important calculations.  These include: (1)~strong-gravity computations of the recoil speed of a black hole binary after merger, for different mass ratios and spins, (2)~comprehensive numerical simulations of three-body interactions with high mass ratios, to represent intermediate-mass black holes, and (3)~detailed numerical analysis of binary-binary encounters and the role  of the Kozai resonance, among others.  Whatever the results of this work, there will be significant new understanding gained on many fronts,  including the information that can be obtained from analysis of the waveforms of the gravitational radiation produced by black hole mergers in dense clusters."
    },
    "0206/astro-ph0206297_arXiv.txt": {
        "abstract": "We present a 63\\,ks {\\it FUSE} observation of the Narrow-Line Seyfert 1 galaxy \\objectname[MGC +05-53-012]{Arakelian~564}. The spectrum is dominated by the strong emission in the \\ion{O}{6}\\,$\\lambda\\lambda1032, 1038$ resonance doublet. Strong, heavily saturated absorption troughs due to Lyman series of Hydrogen, \\ion{O}{6} and \\ion{C}{3} $\\lambda 977$ at velocities near the systemic redshift of Arakelian~564 are also observed. We used the column densities of \\ion{O}{6} and \\ion{C}{3} in conjunction with the published column densities of species observed in the UV and X-ray bands to derive constraints on the physical parameters of the absorber through photoionization modeling. The available data suggest that the UV and X-ray absorbers in Arakelian~564 are physically related, and possibly identical. The combination of constraints indicates that the absorber is characterized by a narrow range in total column density $N_{\\rm H}$ and $U$, centered at $\\log N_{\\rm H} \\approx 21$ and $\\log U \\approx -1.5$, and may be spatially extended along the line of sight. ",
        "introduction": "%  More than half of the Seyfert 1 population shows optical/UV intrinsic absorption associated with their active nucleus \\citep[][ and references therein]{Crenshawea99}. The strong UV absorption lines, Ly$\\alpha$, \\ion{C}{4}, \\ion{N}{5} (and less frequently \\ion{Si}{4} and \\ion{Mg}{2}) are found to be blueshifted, or at rest, with respect to the narrow emission lines, providing an important indication of the presence of a net radial outflow of the absorbing gas. A similar percentage also shows an associated ionized (``warm'') X-ray absorber \\citep{Georgeea98,Reynolds97} that is characterized by high ionization, $U=0.1$--$10$ ($U=Q / 4\\,\\pi \\, r^2 \\, n_{\\rm H}\\,\\,c$, where $Q$ is the number of ionizing photons) and high total Hydrogen column density, $N_{\\rm H} = 10^{21}$--$10^{23}$\\,cm$^{-2}$, which signature is typically the presence of \\ion{O}{7} and \\ion{O}{8} edges. During the last decade evidence has been accumulated indicating that the same gas is responsible for the absorption in the UV and X-ray spectra of Seyfert 1s \\citep{Mathur94,MEWF94,MEW95,Crenshawea99,Krissea00,Monierea01,Brothertonea02}. Although it is not always possible to model the X-ray/UV absorber as a single zone (especially when the complex UV absorption is resolved in multiple velocity components that are characterized by a large range of column densities and ionization) common characteristics of these absorbers have emerged, i.e., they are composed of high ionization, low density, high column density gas that is outflowing and is located in or outside the broad emission-line region (BELR). It is therefore worthwhile to investigate the nature of this nuclear component in active galactic nuclei (AGN) that represents an outflow (or wind) that can carry away a significant amount of kinetic energy at a mass-loss rate comparable to the accretion rate required to fuel the AGN \\citep{MEW95}.  Arakelian~564 (Ark~564, IRAS 22403+2927, MGC +05-53-012) is a bright, nearby Narrow-Line Seyfert 1 (NLS1) galaxy, with $z = 0.02467$ and  $V = 14.6$ mag \\citep{rc3.9catalogue}, and $L_{\\mbox{\\scriptsize 2--10{} keV}} = 2.4 \\times 10^{43}$ \\ergsec{} \\citep[][hereafter Paper I]{Akn564I}. It was the object of an intense multiwavelength monitoring campaign that included simultaneous observations from {\\it ASCA} (2000 June 1 to July 6, Paper I; \\citealt{Poundsea01,Edelsonea02}), {\\it HST} (2000 May 9 to July 8, \\citealt{Collierea01}, Paper II; \\citealt{Crenshawea02}, Paper IV) and from many ground-based observatories as part of an AGN Watch\\footnote{\\anchor{http://www.astronomy.ohio-state.edu/~agnwatch} {All publicly available data and complete references to published AGN Watch papers can be found at http://www.astronomy.ohio-state.edu/$\\sim$agnwatch.}} project \\citep[1998 Nov to 2001 Jan, ][ Paper III]{Shemmerea01}. Akn~564 has shown a strong associated UV absorber (\\citealt{Crenshawea99}, Paper II, Paper IV). There are indications that it also possesses a warm X-ray absorber, as seen by the absorption lines of \\ion{O}{7} and \\ion{O}{8} detected in a {\\it Chandra} spectrum \\citep{Matsumotoea01}\\footnote{http://www.pha.jhu.edu/groups/astro/workshop2001/papers/.}.  In this paper we present the results from a $63$\\,ks {\\it FUSE} observation of Akn~564 obtained on 2001 June 29--30 UT, focusing in particular on the \\ion{O}{6} intrinsic absorption; we investigate the physical properties of the UV and X-ray absorbing gas using the constraints on column densities obtained during the multiwavelength observations of this AGN. In \\S\\ref{fuseobs} we present the data. In \\S\\ref{anal} we describe our analysis methods. In \\S\\ref{photoionization} we test the hypothesis that the Warm UV-X-ray absorber are one and the same through photoionization calculations. In \\S\\ref{discuss} we discuss some implications of our investigation. Our results are summarized in \\S\\ref{summary}. In a forthcoming paper (Romano et al., in preparation) we will analyze the intrinsic SED of Ark~564 and the properties of the gas responsible for the broad emission lines. ",
        "conclusions": "}  The UV absorber in Ark~564 is in a general state of outflow with respect to the systemic redshift (see Table~\\ref{constraints}). A very good agreement is found between the values of the velocity centroids we derive for the species observed in the {\\it FUSE} spectrum and the ones derived for the {\\it HST}/STIS spectrum (Paper IV); therefore, we adopt as the best estimate of the net radial velocity of the UV absorber the value obtained in Paper IV for \\ion{Si}{3} and \\ion{Si}{4}, the least saturated lines: $V_{\\rm out} = -194 \\pm 5$\\,\\kms. The absorption troughs also show the presence of gas which is redshifted with respect to the systemic velocity. This can be explained in part as a saturation effect, as is shown in Figure~3 of Paper IV. Alternatively, a model with more than one kinematic component is required to explain the observed absorption troughs \\citep[i.e.][]{Elvis2000}; in this scenario, in addition to the blueshifted absorption from an outflowing wind, we would be observing redshifted absorption from infalling material, such as an accretion flow. In addition to the continuum source, the absorbing gas must cover a substantial portion of the BELR, since the absorption troughs are much deeper than the continuum level. Assuming the identity of the UV and X-ray absorbing gas, we have used the column densities of the observed species to constrain the physical conditions of the absorber. For the most realistic SED (SED2), we obtained \\centerline{\\vspace{-1.cm}\\hspace{+1.truecm} \\includegraphics[width=10.cm,height=10.0cm]{figure9.ps} } \\figcaption{Same as Fig.~\\ref{tableagn}, but with a ionizing continuum described in \\citet{Laorea97a} and \\citet{Zhengea97} (ZL in Figure~\\ref{allseds}). Solar abundances are assumed. Most constraints are met in $\\log U = [-1.97, -1.54]$ and $\\log N_{\\rm H} = [20.00, 20.79]$ (see \\S\\ref{absorber}). \\label{lzsed}}  \\vspace{0.6cm}  \\noindent $\\log N_{\\rm H} = [19.95, 21.27]$ and $\\log U = [-1.86, -1.02]$. These constraints can be used to determine the size of the absorber, its distance from the central continuum source, and the mass outflow rate. For the following order-of-magnitude arguments we adopt the mean values $U = 0.0363$ and $N_{\\rm H}=4.07 \\times 10^{20}$\\,cm$^{-2}$. The size of the absorber is  derived from the total column density, $r_{\\rm abs} = 4.07 \\times 10^{20} \\,n_{\\rm H}^{-1}$\\, cm. For SED2 the number of ionizing photons is $Q = 6.68 \\times 10^{55}$\\,s$^{-1}$, so the distance from the continuum source is $R_{\\rm abs} =  7.00 \\times 10^{22} \\, n_{\\rm H}^{-1/2}$\\,cm. Using the lower limit on $R_{\\rm abs} > 95$\\, pc found in Paper IV, this would indicate a total density $n_{\\rm H} > 5.70 \\times 10^{4}$\\,cm$^{-3}$. Assuming uniform density, and considering that $r_{\\rm abs} \\ll R_{\\rm abs}$, the mass of the outflowing gas is $M_{\\rm abs} = 2.11 \\times 10^{10}\\,\\,f\\,n_{\\rm H}^{-1}\\,\\Msun$, where $f$ is the covering factor, i.e., the fraction of the sky covered by the absorber as seen at the central source. The mass outflow rate is then $\\dot{M}_{\\rm abs} = M_{\\rm abs}\\,V_{out}\\,\\, / r_{\\rm abs} = 3.17 \\times 10^{4}\\,\\,f$\\,\\,\\Msun\\,yr$^{-1}$ and the outflow carries out a rate of kinetic energy $\\dot{M}_{\\rm abs} \\,V_{\\rm out}^2 / 2 = 3.76 \\times 10^{44}\\,f\\,\\,$\\ergsec. To power Ark~564 at the observed luminosity ($L_{\\rm bol} = 10 \\times L_{\\mbox{\\scriptsize 2--10{} keV}} = 2.4 \\times 10^{44}$\\,\\ergsec) at an efficiency $\\eta = 0.1$, an accretion rate $\\dot{M}_{\\rm acc} = 1.8 \\times 10^{-3} \\left( L_{44} / \\eta\\right) = 4.3 \\times 10^{-2}\\,\\Msun\\,$yr$^{-1}$ is required ($L_{44}$ is the bolometric luminosity in units of 10$^{44}$\\,\\ergsec).  The outflow carries out a kinetic luminosity about one order of magnitude smaller than the observed radiative luminosity of the source. However, the mass outflow rate is uncomfortably large unless the covering factor is very small. If $\\dot{M}_{\\rm abs} \\la \\dot{M}_{\\rm acc}$, then it implies $f \\la 10^{-6}$. Alternatively, our assumption $r_{\\rm abs} \\ll R_{\\rm abs}$ might not be valid. The absorber might be an extended, low density region. The assumption of a uniform density gas may not be strictly valid and the ionization parameter and density that we deduced should only be considered as ``average'' values. Recent {\\it Chandra} observations have found extended warm gas in some AGN \\citep{Sakoea00} with physical characteristics similar to that of a warm absorber, but seen in emission. So it is quite likely that the warm absorber in Ark~564 is also spatially extended along the line of sight.  In Paper IV the UV absorber was modeled as a single zone with quasi-solar abundances (Carbon depletion being the main departure) and the  best-fit values of  $\\log U = -1.48 $ and $\\log N_{\\rm H} = 21.21$  are consistent with our limits. \\citet[][Paper IV]{Crenshawea02} over-predicted Carbon and Oxygen column densities: $N_{\\rm C\\, III} = 5.2 \\times 10^{15}$ cm$^{-2}$, $N_{\\rm O\\,VI} = 2.4 \\times 10^{17}$ cm$^{-2}$ (cf.\\ our measurements: $N_{\\rm C\\, III} = 3.2 \\times 10^{14}$ cm$^{-2}$ and $N_{\\rm O\\,VI} = 5.7 \\times 10^{15}$ cm$^{-2}$). These predictions, however, are consistent with the upper-end values of our range of parameter space. We note that our modeling did not require Carbon to be depleted.  Finally, we can compare our solutions of $\\log N_{\\rm H} = 21 $ and $\\log U = -1.5$, with the preliminary results of \\citet{Matsumotoea01} based on analysis of a 50\\,ks {\\it Chandra} observation of Ark~564. Their curve of growth analysis on the absorption lines indicates that $N_{\\rm O\\, VII}  = 3.2 \\times 10^{17}$ cm$^{-2}$, $N_{\\rm O\\, VIII} = 1   \\times 10^{18}$ cm$^{-2}$, $N_{\\rm Ne\\,IX}   = 3.2 \\times 10^{17}$ cm$^{-2}$, and $N_{\\rm Ne\\, X}   = 1   \\times 10^{17}$ cm$^{-2}$, suggestive of $\\log N_{\\rm H} = 21$ and $\\log \\xi = 1.6$--2. While there is agreement between the values of $\\log N_{\\rm H}$ and $N_{\\rm O\\, VII}$, the column densities they measure for \\ion{O}{7} do not agree with the ones derived in Paper IV from the upper limits on the bound-free optical depths in the {\\it ASCA} spectrum (Paper I). Given the high $N_{\\rm O\\, VIII}$ we would expect an edge would be observable in the {\\it Chandra} spectrum."
    },
    "0206/astro-ph0206068_arXiv.txt": {
        "abstract": "The angular correlation function of the background shear-foreground galaxy distribution probes the three dimensional cross power spectrum between mass and galaxies. The same cross power spectrum is also probed when foreground galaxy distribution  is cross-correlated with a distribution of background sources disjoint in redshift space. The kernels  that project three dimensional clustering to the two dimensional angular space is different for these two probes. When combined, they  allow a study of the galaxy-mass cross power spectrum from linear to non-linear scales.  By inverting the background shear-foreground galaxy correlation function measured by the Sloan Digital Sky Survey, we present a first estimate of the cross power spectrum between mass and galaxies at low redshifts. ",
        "introduction": "The current and upcoming wide field imaging data, such as from the Sloan Digital Sky Survey, allow detailed studies on the angular clustering of galaxies and quasars (see, e.g., Dodelson et al. 2001; Scranton 2002). In addition to such studies, these data also allow investigations related to the cross clustering between galaxies and mass.  With the so-called galaxy-galaxy lensing (Blandford et al. 1991; Bartelmann \\& Schneider 2002 for a recent review), one measures the correlation between foreground galaxies and the background shear surrounding these galaxies. This results in a measuremnt of  the correlation between foreground galaxies and the dark matter distribution traced by these galaxies. Under the assumption that a single galaxy reside in each  dark matter halo, one can use the observed correlation function to constrain some physical properties of halos  (Fischer et al. 2000). Alternatively, more detailed, and increasingly complicated, models can also be introduced such that one takes in to account the fact that more than one galaxy may be present in dark matter halos at the high end of the mass function (Guzik \\& Seljak 2001, 2002).  In addition to the background shear-foreground galaxy correlation function, the cross-correlation function between  a sample of foreground galaxies and background sources, such as quasars, also probe the same galaxy-dark matter cross power spectrum (Moessner et al. 1997). This correlation results from the fact that number counts of background galaxies are affected by lensing magnification via the intervening dark matter distribution traced by foreground galaxies.  In this {\\it letter}, we will briefly consider complimentary properties of these two lensing probes of the same galaxy-mass cross power spectrum. The foreground galaxy-background source correlation and the foreground galaxy-background shear correlation  have unique properties in that they probe two different regimes of the galaxy-dark matter cross-power spectrum. We will describe this difference in terms of the kernel involved with the projection of three dimensional clustering to the two dimensional angular space. Finally,  we will estimate the cross power spectrum between galaxies and mass by inverting the published galaxy-shear correlation function from the Sloan Digital Sky Survey. When necessary, we will illustrate our results using the currently favored $\\Lambda$CDM cosmology with $\\Omega_m=0.35$ and $\\Omega_\\Lambda=0.65$ and use inputs, such as the redshift distributions necessary for inversions, from published results in the literature. ",
        "conclusions": "We have discussed two probes of the cross power spectrum between galaxies and mass which effectively use two aspects of weak lensing involving shearing of background images in one case and the magnification of background images in the other. The two probes, though similar in most aspects, probe different physical scales of the galaxy-mass cross power spectrum due to a subtle difference in the kernel which projects three dimensional clustering to the observable two-dimensional angular space. When combined, effectively with the same sample of foreground galaxies whose redshift distribution is known a priori, these two probes  allow a study of the galaxy-mass cross power spectrum from linear to non-linear scales.  While the galaxy-mass cross power spectrum provides information on how galaxies are correlated with mass, the galaxy-galaxy power spectrum provides information on biasing. Thus, a complete study, which can be easily carried out with imaging data such as from the Sloan survey is to measure the angular correlation of foreground galaxies, and associated lensing correlation functions discussed here. The three functions can then be inverted in a consistent manner to obtain much needed knowledge on galaxy clustering relative to mass. As a first example of such an approach, we have inverted the background shear-foreground galaxy correlation function measured by the Sloan Digital Sky Survey and have provided a first estimate of the cross power spectrum between mass and galaxies at low redshifts as appropriate for Sloan sample of galaxies in the r' band."
    },
    "0206/hep-ph0206131_arXiv.txt": {
        "abstract": "We present the results of a detailed study of how isocurvature axion fluctuations are converted into adiabatic metric perturbations through axion decay, and discuss the constraints on the parameters of  pre-big bang cosmology needed for consistency with present CMB-anisotropy data. The large-scale normalization of temperature fluctuations has a non-trivial dependence both on the mass and on the initial value of the axion. In the simplest, minimal models of pre-big bang inflation, consistency with the COBE normalization requires a slightly tilted (blue) spectrum, while a strictly  scale-invariant spectrum requires mild modifications of the minimal  backgrounds at large curvature and/or string coupling. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206224_arXiv.txt": {
        "abstract": "Brane world cosmologies seem to provide an alternative explanation for the present accelerated stage of the Universe with no need to invoke either a cosmological constant or an exotic \\emph{quintessence} component. In this paper we investigate statistical properties of gravitational lenses for some particular scenarios based on this large scale modification of gravity. We show that a large class of such models are compatible with the current lensing data for values of the matter density parameter $\\Omega_{\\rm{m}} \\leq 0.94$ ($1\\sigma$). If one fixes $\\Omega_{\\rm{m}}$ to be $\\simeq 0.3$, as suggested by most of the dynamical estimates of the quantity of matter in the Universe, the predicted number of lensed quasars requires a slightly open universe with a crossover distance between the 4 and 5-dimensional gravities of the order of $1.76 H_o^{-1}$. ",
        "introduction": "The results of observational cosmology in the last years have opened up an unprecedented opportunity to test the veracity of a number of cosmological scenarios as well as to establish a more solid connection between particle physics and cosmology. The most remarkable finding among these results comes from distance measurements of type Ia supernovae (SNe Ia) that suggest that the expansion of the Universe is speeding up, not slowing down \\cite{perlmutter}. As widely known such a result poses a crucial problem for all CDM models since their generic prediction is a decelerating universe ($q_{o} > 0$), whatever the sign adopted for the curvature parameter. Indirectly, similar results have also been obtained, independent of the SNe Ia analyses, by combining the latest galaxy clustering data with CMB measurements \\cite{efs}.  To reconcile these observational results with theory, cosmologists have proposed more general models containing a negative-pressure dark component that would be responsible for the present accelerated stage of the Universe. Although a large number of pieces of observational evidence have consistently suggested a universe composed of $\\sim 2/3$ of dark energy, the exact nature of this new component is not well understood at present. Among the several candidates for dark energy discussed in the recent literature, the simplest and most theoretically appealing possibility is the vacuum energy or cosmological constant. Despite the serious problem that arises when one considers a nonzero vacuum energy \\cite{wein}, models with a relic cosmological constant ($\\Lambda$CDM) seem to be our best description of the observed universe, being considered as a serious candidate for standard cosmology.  On the other hand, motivated by particle physics considerations, there has been growing interest in cosmological models based on the framework of brane-induced gravity \\cite{ark,dvali,deff,randall}. The general principle behind such models is that our 4-dimensional Universe would be a surface or a brane embedded into a higher dimensional bulk space-time on which gravity can propagate. In some of these scenarios, there is a certain crossover scale $r_c$ that defines what kind of gravity an observer on the brane will observe. For distances shorter than $r_c$, such an observer will measure the usual 4-dimensional gravitational $1/r^{2}$ force whereas for distances larger than $r_c$ the  gravitational force follows the 5-dimensional $1/r^{3}$ behavior. In this way, gravity gets weaker at cosmic distances and, therefore, it is natural to think that such an effect has some implications on the dynamics of the Universe \\cite{deff1}.  Several aspects of brane world cosmologies have been explored in the recent literature. For example, the issue related to the cosmological constant problem has been addressed \\cite{ccp} as well as evolution of cosmological perturbations in the gauge-invariant formalism \\cite{brand}, cosmological phase transitions \\cite{cpt}, inflationary solutions \\cite{cpt1}, baryogenesis \\cite{dvali99}, stochastic background of gravitational waves \\cite{hogan1}, singularity, homogeneity, flatness and entropy problems \\cite{aaa},  among others (see \\cite{hogan} for a discussion on the different perspectives of brane world models). From the observational viewpoint, however, the present situation is somewhat controversial. While the authors of Refs. \\cite{deffZ,dnew} have shown that such models are in agreement with the most recent cosmological observations (for example, they found that a flat universe with $\\Omega_{\\rm{m}} = 0.3$ and $r_c \\simeq 1.4H_o^{-1}$ is consistent with the currently SNe Ia + CMB data), the authors of Ref. \\cite{avelino} have claimed that a larger sample of SNe Ia data can also be used to rule out these models at least at the 2$\\sigma$ level. Recently, one of us \\cite{alcaniz} used measurements of the angular size of high-$z$ compact radio sources to show that the best fit model for these data is a slightly closed universe with $\\Omega_{\\rm{m}} \\simeq 0.06$ and a crossover radius of the order of 0.94$H_o^{-1}$.   For the reasons presented earlier, the comparison between any alternative cosmology and $\\Lambda$CDM models is very important. In this concern, statistical properties of gravitational lenses may be an interesting tool because, as is well known, they provide restrictive limits on the vaccum energy contribution (see, for instance, \\cite{1CSK}). On the other hand, in brane world models the distance to an object at a given redshift $z$ is smaller than the distance to the same object in $\\Lambda$CDM models (assuming the same value of $\\Omega_{\\rm{m}}$). Therefore, we expect that the constraints coming from lensing statistics will be weaker for these models than for their $\\Lambda$CDM counterparts.  In this paper, we explore the implications of gravitationally lensed QSOs for models based on the framework of the brane-induced gravity of Dvali {\\it et al.} \\cite{dvali} that have been recently proposed in Refs. \\cite{deff,deff1}. We restrict our analysis to accelerated models, or equivalently, models that have a ``self-inflationary\" solution with $H \\sim r_c^{-1}$ ($H$ is the Hubble parameter). As explained in \\cite{deff}, in such scenarios, the bulk gravity sees its own curvature term on the brane as a negative-pressure dark component and accelerates the Universe.  This paper is organized as follows. In section 2, we present the basic field equations and distance formulas relevant for our analysis. We then proceed to analyze the constraints from lensing statistics on these models in section 3. In section 4 our main conclusions are presented. ",
        "conclusions": "The recent observational evidences for a presently accelerated stage of the Universe have stimulated renewed interest for alternative cosmologies. In general, such models contain an unkown negative-pressure dark component that explains the SNe Ia results and reconciles the inflationary flatness prediction ($\\Omega_{\\rm{T}} = 1$) with the dynamical estimates of the quantity of matter in the Universe ($\\Omega_{\\rm{m}} \\simeq 0.3 \\pm 0.1$). In this paper we have focused our attention on another dark energy candidate, one arising from gravitational \\emph{leakage} into extra dimensions \\cite{deff,deff1}. We have shown that some particular scenarios based on this large scale modification of gravity are in agreement with the current gravitational lensing data for values of $\\Omega_{\\rm{m}} \\leq 0.93$ ($1\\sigma$). If one fixes $\\Omega_{\\rm{m}}$ to be $\\simeq 0.3$, the predicted number of lensed quasars requires $\\Omega_{r_{c}} \\simeq 0.08$. This is a slightly open universe with a crossover radius of the order of $r_c \\simeq 1.76 H_o^{-1}$."
    },
    "0206/astro-ph0206012_arXiv.txt": {
        "abstract": "We report the results of our continued study of arcminute scale anisotropy in the Cosmic Microwave Background (CMB) with the Berkeley-Illinois-Maryland Association (BIMA) array. The survey consists of ten independent fields selected for low infrared dust emission and lack of bright radio point sources. With observations from the Very Large Array (VLA) at $4.8$ GHz, we have identified point sources which could act as contaminants in estimates of the CMB power spectrum and removed them in the analysis. Modeling the observed power spectrum with a single flat band power with average multipole of $\\ell_{eff} = 6864$, we find $\\Delta T=14.2^{+4.8}_{-6.0}\\,\\mu$K at $68\\%$ confidence. The signal in the visibility data exceeds the expected contribution from instrumental noise with $96.5\\%$ confidence. We have also divided the data into two bins corresponding to different spatial resolutions in the power spectrum. We find $\\Delta T_1=16.6^{+5.3}_{-5.9}\\,\\mu$K at $68\\%$ confidence for CMB flat band power described by an average multipole of $\\ell_{eff} = 5237$ and $\\Delta T_2<26.5\\,\\mu$K at $95\\%$ confidence for $\\ell_{eff} = 8748$. ",
        "introduction": "\\label{sec:intro} Fluctuations in the distribution of matter at the epoch of recombination create large angular scale anisotropy in the Cosmic Microwave Background (CMB). This primordial anisotropy has been studied extensively at degree and sub-degree angular scales in order to place constraints on the parameters of cosmological models (Halverson et al. 2002, de Bernardis et al. 2002, Scott et al. 2002, Lee et al. 2001, Miller \\ea 2002, Padin et al. 2001). At arcminute scales, the primordial anisotropy is damped to negligible amplitude due to photon diffusion and the finite thickness of the last scattering surface (Hu \\& White 1997). On these smaller scales, secondary anisotropies such as the Sunyaev-Zeldovich effect (SZE) are expected to dominate the signal of the CMB power spectrum (Haiman $\\&$ Knox 1999). Studies of secondary anisotropy in the CMB have the potential to be a powerful probe of the growth of structure in the Universe.  In this paper, we report results from an ongoing program using the Berkeley-Illinois-Maryland Association (BIMA) interferometer to search for arcminute-scale CMB anisotropy. Discussion of the instrument, data reduction, expected signals (from both primary and secondary anisotropies) and previous measurements is included in earlier publications (Holzapfel et al. 2000 and Dawson et al. 2001, hereafter H2000 and D2001 respectively). We describe observations and criteria for field selection in \\S\\ref{sec:obs}. The Bayesian likelihood analysis used to constrain the CMB anisotropy is described in \\S\\ref{sec:anal}. The results are presented in \\S\\ref{sec:results} including a discussion of tests for systematic errors in the analysis. Finally, in \\S\\ref{sec:con}, we present the conclusion and comparison to simulations of the SZE. ",
        "conclusions": "\\label{sec:con} Over the course of three summers, we have used the BIMA array in a compact configuration at $28.5\\,$GHz to search for CMB anisotropy in ten independent $6.6^\\prime$ FWHM fields. With these observations, we have detected arcminute scale anisotropy at better than $95\\%$ confidence. In the context of an assumed flat band power model for the CMB power spectrum, we find $\\Delta T=14.2^{+4.8}_{-6.0}\\,\\mu$K at $68.3\\%$ confidence with sensitivity on scales that correspond to an average harmonic multipole $\\ell_{eff} = 6864$. We also present results after dividing the visibility data into two bins of different spatial resolution. We find $\\Delta T_1=16.6^{+5.3}_{-5.9}\\,\\mu$K at $68.3\\%$ confidence on scales corresponding to an average harmonic multipole $\\ell_{eff} = 5237$ and $\\Delta T_2 < 26.5 \\,\\mu$K at $95\\%$ confidence at $\\ell_{eff} = 8748$. The results of the VLA observations appear to exclude the possibility of point source contamination and there is no indication of an obvious systematic error that would bias the observations.  Mason et al. (2002) have also reported a detection of excess power at somewhat larger angular scales. They find $\\Delta T=22.5^{+2.5}_{-3.6}\\,\\mu$K for data in the range $2010 \\, < \\, \\ell \\, < 4000$. Although this measurement is at a lower $\\ell$ than the BIMA results, it is significantly higher than the expected power due to primordial anisotropy.  If the signal observed in the BIMA fields is indeed caused by CMB anisotropy, there are a host of possible sources for excess power such as primary anisotropy, thermal SZE, kinetic SZE, patchy reionization, and the Ostriker-Vishniac effect. Of these candidates for CMB anisotropy, the thermal SZE from clusters of galaxies is expected to dominate on the scales where the BIMA instrument is most sensitive (see for example, Gnedin \\& Jaffe 2001). Analytic models and simulations of cluster formation predict $\\Delta T$ values that range from $4.3\\,\\mu$K to $15.0\\,\\mu$K on angular scales of approximately two arcminutes. Figure 7 compares the results of this paper to the theoretical models.  The non-Gaussian characteristics of the CMB power spectrum caused by the thermal SZE may increase the uncertainty in measurements of the power spectrum due to sample variance. Current models suggest that these effects increase the uncertainty by a factor of 3 over what is expected for the sample variance of a Gaussian distributed signal at $\\ell \\sim 5000$ (White, Hernquist, \\& Springel 2002). Based on this argument, the effect of non-Gaussian sample variance contributes a standard deviation of $3 \\times \\sqrt{2/N} \\,  \\Delta T = 6.0 \\, \\mu$K to the anisotropy measurements, where $N\\sim 100$ is the number of independent pixels in the $u$-$v$ range $0.63-1.1\\,\\kl$ and $\\sqrt{2/N} \\, \\Delta T$ is the sample variance in a Gaussian distributed signal. The uncertainty due to sample variance is approximately equal to the statistical uncertainty reported in this paper, increasing the overall uncertainty by $40 \\%$, in agreement with the predictions of Zhang, Pen $\\&$ Wang (2002). The predicted effect of sample variance on the measurements is represented by the extended error bars in Figure 7.   \\vskip 20 pt  We thank the entire staff of the BIMA observatory for their many contributions to this project, in particular Rick Forster and Dick Plambeck for their assistance with both the instrumentation and observations. Nils Halverson and Martin White are thanked for stimulating discussions concerning data analysis. We are grateful for the scheduling of time at the VLA in support of this project that has proved essential to the point source treatment. This work is supported in part by NASA LTSA grant number NAG5-7986, NSF grant 0096913, and the David and Lucile Packard Foundation. The BIMA millimeter array is supported by NSF grant AST 96-13998. A.M. is supported by Hubble Grant ASTR/HST-HF-0113.  \\newpage \\markright{REFERENCES}"
    },
    "0206/astro-ph0206362_arXiv.txt": {
        "abstract": "Through a 3d-modeling of ASCA observations, we performed a spatially resolved X-ray spectroscopic study, extending to radii exceeding 150 kpc, for a sample of 9 groups of galaxies. Combined with published ROSAT results, we conclude that these systems generally exhibit a strong temperature decline at outer radii. In our best case, NGC3268, this corresponds to a flattening of the entropy profile at a level of $\\sim400$ keV cm$^2$. This value is high compared both to the observed entropy floor of $\\sim100$ keV cm$^2$ and to the expected value from gravitational heating. We suggest that the observed entropy profile in most groups at densities exceeding 500 times the critical is purely driven by non-gravitational heating processes. After comparison with a larger sample of groups and clusters, we conclude that there is a variation in the level of non-gravitational heating between $\\sim100$ keV cm$^2$ and $\\sim400$ keV cm$^2$ within every system. Using models of cluster formation as a reference frame, we established that the accreted gas reaches an entropy level of $400$ keV cm$^2$ by redshift $2.0-2.5$, while such high entropies where not present at redshifts higher than $2.8-3.5$, favoring nearly instantaneous preheating. Adopting galactic winds as a source of preheating, and scaling the released energy by the observed metal abundance, the variation in the preheating could be ascribed mostly to variation in the typical overdensity of the energy injection, $\\sim30$ for an entropy floor ($100$ keV cm$^2$) and to $\\sim5$ for an entropy of $400$ keV cm$^2$. ",
        "introduction": "Comparative studies of the scaling relations in clusters of galaxies reveal strong deviations of the observed relations from predictions based on self-similar collapse (Evrard \\& Henry 1991; Bower 1996; Loewenstein 2000; Finoguenov, Reiprich, B\\\"ohringer 2001). These deviations are best characterized by an entropy floor in the X-ray emitting gas. While negligible compared to the entropy due to the accretion shock in large clusters, such preheating\\footnote{We define preheating as a change in the initial adiabat for the gas before its accretion into the potentials of groups and clusters. The plausible sources of preheating include galactic winds and AGN activity, as well as gravitational shocks associated with the formation of large scale structure.} leads to very extended, low-density gas distributions on the scales of groups, that cause a steepening of most scaling relations (Ponman, Cannon, Navarro 1999, hereafter PCN). When the temperature of the preheated gas reaches the level comparable to the virial temperature, simulations predict adiabatic accretion of the gas (\\eg\\ Tozzi, Scharf, Norman 2000). This also implies strong temperature gradients, proportional to the gas density to the 2/3 power ($\\gamma=5/3$). As noted by Loewenstein (2000), the observation of this phenomena is a critical test for the importance of preheating, as well as other characteristics of cluster collapse. So far no clear examples of systems with an adiabatic gas distribution have been found.   Throughout this {\\it Paper} we will define the entropy as $S={kT_e / n_e^{2/3}}h^{-1/3}$ in keV cm$^2$, following PCN. The observational data is presented assuming $H_{\\rm o}=50$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0206/astro-ph0206154_arXiv.txt": {
        "abstract": "\\noindent  We provide a simple theoretical model for the quasar luminosity function at high redshifts that naturally reproduces the statistical properties of the luminous SDSS quasar sample at redshifts $z\\sim4.3$ and $z\\ga 5.7$. Our model is based on the assumptions that quasar emission is triggered by galaxy mergers, and that the black hole mass is proportional to a power-law in the circular velocity of the host galactic halo, $v_c$. We assume that quasars shine at their Eddington luminosity over a time proportional to the mass ratio between the small and final galaxies in the merger.  This simple model fits the quasar luminosity function at $z\\sim2$--$3$, reproduces the normalization and logarithmic slope ($\\beta\\sim-2.58$) at $z\\sim4.3$, explains the space density of bright SDSS quasars at $z\\sim6.0$, reproduces the black hole -- halo mass relation for dormant black holes in the local universe, and matches the estimated duty cycle of quasar activity ($\\sim 10^7$ years) in Lyman-break galaxies at $z\\sim 3$.  Based on the derived luminosity function we predict the resulting gravitational lensing rates for high redshift quasars.  The lens fractions in the SDSS samples are predicted to be $\\sim 2\\%$ at $z\\sim4.3$ and $\\sim 10\\%$ at $z\\ga5.7$.  Interestingly, the limiting quasar luminosity in our best-fit relation $L \\propto v_c^5/G$, scales as the binding energy of the host galaxy divided by its dynamical time, implying that feedback is the mechanism that regulates black hole growth in galactic potential wells. ",
        "introduction": "While the quasar luminosity function has been studied extensively at redshifts below $z\\sim3$ (e.g. Boyle, Shanks \\& Peterson~1988; Hartwick \\& Schade 1990; Pei~1995), the Sloan Digital Sky Survey (SDSS; Fukugita et al.~1996; Gunn et al.~1998; York et al.~2001) has in recent years substantially increased the number of quasars known at $z\\ga3.5$ (Fan, Strauss et al.~2001a,b; Schneider et al.~2001). Two samples of very high redshift SDSS quasars have been presented to date. The first of these is a sample of 39 luminous quasars with redshifts in the range $3.6<z<5.0$ and a median of 4.3. The second sample consists of 4 quasars at $z\\ga5.7$, including the quasar with the highest known redshift ($z=6.28$).  These samples begin to sketch out the luminosity function of quasars at $z\\ga4$ and are very important for studies of the ionizing background radiation and of quasar evolution around the epoch of reionization (Haiman \\& Loeb~1998). Although the logarithmic slope of the luminosity function for bright quasars admits a universal value of $\\sim-3.5$ at redshifts $z\\la 3$, Fan et al.~(2000a) find a shallower logarithmic slope of $\\sim-2.5$ for their luminous quasars at $z\\ga 3.5$. At first sight this is a puzzling result because the exponential tail of the Press-Schechter~(1974) mass function describing the space density of the massive host galaxies in which bright quasars are thought to reside, becomes progressively steeper with increasing redshift.  There are several published models of varying complexity to describe the observed evolution of the quasar luminosity function within hierarchical structure formation theory (Efstathiou \\& Rees 1988; Small \\& Blandford 1992; Haehnelt \\& Rees 1993; Haiman \\& Loeb~1998; Haehnelt Natarajan \\& Rees~1998; Kauffmann \\& Haehnelt~2000).  In these models, the luminosities and lifetimes of the quasars are governed by the black hole mass and the supply of cold accreting gas. The abundance and evolution of the supermassive black holes assumed to power the quasars, are linked to the evolution of the mass-function of galactic halos.  The density of bright quasars is observed to decline rapidly with redshift below $z\\sim2$ (Boyle, Shanks \\& Peterson~1988; Hartwick \\& Schade 1990). This decline has been explained (Kauffman \\& Haehnelt~2000; see also Cavaliere \\& Vittorini~2000) in terms of decreases in the merger rate and the availability of cold gas to fuel the black holes.  In addition, the density of bright quasars is also observed to decline beyond a redshift of $z\\sim2$ (Warren, Hewett \\& Osmer~1994; Schmidt, Schneider \\& Gunn~1995; Kennefick, Djorgovski \\& de Carvalho~1995). Under the assumption that quasar activity is related to the formation rate of halos, model luminosity functions have been constructed that successfully describe this decline and the shape of the luminosity function between $z\\sim2$ and $z\\sim3.5$ (Efstathiou \\& Rees~1988; Haehnelt \\& Rees~1993; Haiman \\& Loeb~1998; Haehnelt, Natarajan \\& Rees~1999; Haiman, Madau \\& Loeb~1999). These models predict the luminosity function at still higher redshifts. With the recent influx of results on the quasar luminosity at redshifts above $z\\sim4$, it is interesting to revisit the question of modeling the evolution of the quasar luminosity function at high redshifts. In this paper, we demonstrate that a very simple model can explain the new results on the high redshift quasar luminosity function.  Our model for the evolution of the quasar luminosity function is based on the Press-Schechter~(1974) mass function and halo merger rates computed with the excursion set formalism (Bond et al.~1991; Lacy \\& Cole~1993). We show that this model can reproduce all of the known properties of the luminosity function above $z\\sim2$. We begin in \\S\\ref{models} by describing an approximate version of our model (Haiman \\& Loeb~1998; hereafter HL98), followed by the full discussion based on the halo merger rates. In \\S\\ref{lumfuncobs} and \\S\\ref{cp} we discuss the data on the quasar luminosity function and its comparison with our model.  Finally we discuss the implications of our model for gravitational lensing in high-redshift samples of quasars in \\S\\ref{lens}, and present our main conclusions in \\S\\ref{conclusion}. Throughout the paper we assume density parameters of $\\Omega_{m}=0.35$ in matter, $\\Omega_\\Lambda=0.65$ in a cosmological constant, and a Hubble constant of $H_0=65~{\\rm km\\,s^{-1}\\,Mpc^{-1}}$. For calculations of the Press-Schechter~(1974) mass function we assume a primordial power-spectrum with $n=1$ and the fitting formula to the exact transfer function of Bardeen et al.~(1986). ",
        "conclusions": "\\label{conclusion}  We predict the quasar luminosity function under the hypothesis that quasar activity is associated with galaxy mergers.  Galactic halos are assumed to each host a black hole having a mass determined from a power-law relation with halo circular velocity, $M_{\\rm bh}\\propto v_{\\rm c}^\\gamma$. Upon merger we assume that the gas from the accreted halo is driven to the center, and that the quasar then shines at its Eddington rate for a short period of time, proportional to the (smaller than 1/2) mass ratio between the small and final halos in the merger. Our model has 3 free parameters but at least the following 5 constraints: the luminosity function slope and normalization at $z\\sim 2.5$, the evolution of the luminosity function slope and normalization, and the curvature of the luminosity function.  Our model is simplified by the fact that the effect of reionization on the luminosity function is limited to luminosities much fainter than those detectable in current quasar surveys.  This very simple model reproduces all known properties of the luminosity function above $z\\sim2$, including the recently measured luminosity function slope at $z\\sim4.3$ (Fan et al.~2001a) and the space density of quasars at $z\\sim6.0$ (Fan et al.~2001b). The model suggests a black hole -- halo mass relation, $M_{\\rm bh}\\propto v_{\\rm c}^5$ consistent with that found by Ferrarese~(2002) at $z=0$, and a duty cycle of $\\sim10^7$ yr consistent with previous determinations (Yu \\& Tremaine~2002; Steidel et al. 2002). Using the derived luminosity function we predict the resulting lensing rates for the highest redshift samples. The recently published flux limited samples of SDSS quasars at $z\\sim4.3$ and $z\\ga5.7$ are predicted to have multiple image fractions of $\\sim2\\%$ and $\\sim10\\%$, respectively.  The model luminosity function computed using the halo merger rates is consistent with the density of faint quasars at low redshifts, and predicts a decline in the density of bright quasars towards low redshifts, but predicts a luminosity function slope that is shallower than observed. As has been noted previously, a more careful analysis (e.g. Kauffmann \\& Haehnelt~2000) is required to understand the evolution of the quasar luminosity function at low redshifts. Nevertheless, the consistency of the faint quasar normalization stands in contrast to models that associate quasar activity with the halo formation rate, as those predict no faint quasars at low redshifts.  Thus one is led to conclude that mergers are the likely trigger for quasar activity at all redshifts.  The $M_{\\rm bh}$--$v_{\\rm c}$ relation that provides the best fit to the data may also have a simple physical origin. First, consider dimensional analysis. If we make the minimal assumption that the limiting luminosity of a quasar is only a function of the halo circular velocity $v_{\\rm c}$, then the only dimensional parameters in the problem are $v_{\\rm c}$ and $G$ (Newton's constant). The only combination of these parameters that has dimensions of luminosity is $v_{\\rm c}^5/G$ and so we conclude \\begin{equation} L_{\\rm Edd} \\propto {v_{\\rm c}^5\\over G}. \\end{equation} {\\it What is the physics behind this relation?} The value of $v_{\\rm c}^5/G$ amounts to depositing the entire binding energy of a self-gravitating system during its dynamical time (so that it does not have time to adjust). The binding energy of a self-gravitating mass $M$ is $\\sim M v_{\\rm c}^2$ and the dynamical time is $\\sim r/v_{\\rm c}$. The ratio between these quantities is $\\sim M v_{\\rm c}^3/r$. Using the virial relation $GM/r\\sim v_{\\rm c}^2$, we get the energy deposition rate that would unbind a self-gravitating system on its dynamical time, $\\sim v_{\\rm c}^5/G$. A quasar may therefore unbind the gas in the galaxy around it\\footnote{We implicitly assume that the central region of the galaxy (such as the proto-bulge) is characterized by a circular velocity similar to that of the galactic halo, even though it is dominated by baryons. This assumption follows naturally from the nearly flat rotation curves of nearby galaxies.}  if its power output is too large\\footnote{Note that the power output from the quasar may include both radiation and mechanical energy. Substantial outflows are inferred to exist in broad absorption line quasars and radio galaxies [see recent summary in Furlanetto \\& Loeb (2002), and references therein].}  (Silk \\& Rees 1998; Ciotti \\& Ostriker~2001). As the mass of the black hole increases and the quasar's Eddington luminosity approaches this limit, the feedback will generate a powerful galactic wind and terminate the accretion that feeds the quasar.  Hence, the quasar phenomenon may be self terminating (similarly to the formation of a proto-star).  Of course, if some of the quasar energy escapes (due to incomplete absorption, partial covering factor or efficient cooling by the surrounding gas), then the feedback-limited quasar luminosity would be higher than $v_c^5/G$.  As long as the feedback from all quasars encounters self-similar conditions in different halos, we may write \\begin{equation} \\label{led} L_{\\rm Edd}= \\beta_\\epsilon {v_{\\rm c}^5\\over G} , \\label{eq:beta} \\end{equation} where $\\beta_\\epsilon$ is a constant that is related to the inferred value of $\\epsilon_{\\rm o}$.  Equations (\\ref{vc}), (\\ref{eps}), (\\ref{led}) and the best-fit value of $\\epsilon_{\\rm o}\\sim 10^{-5.2}$ together with the relation $L_{\\rm Edd}=1.4\\times 10^{38}(M_{\\rm bh}/M_\\odot)~{\\rm erg~s^{-1}}$, yield $\\beta_\\epsilon\\approx 50$. We have shown that the assumption of a constant $\\beta_\\epsilon$ in equation~(\\ref{eq:beta}) leads to a luminosity function that describes the data over a wide range of redshifts. The constancy of $\\beta_\\epsilon$ (which is also expected from dimensional analysis) implies that feedback may indeed be the mechanism that regulates the growth of supermassive black holes in galactic potential wells.  The similarity between our deduced values of $\\gamma=5$ and $\\epsilon_{\\rm o}= 10^{-5.2}$ at high redshifts and the values inferred for remnant black holes in the local universe (Ferrarese 2002; Tremaine et al. 2002), provides a strong testimony to this effect."
    },
    "0206/hep-ph0206196_arXiv.txt": {
        "abstract": "We calculate the equation of state for degenerate quark matter to leading order in hard-dense-loop (HDL) perturbation theory. We solve the Tolman-Oppenheimer-Volkov equations to obtain the mass-radius relation for dense quark stars. Both the perturbative QCD and the HDL equations of state have a large variation with respect to the renormalization scale for $\\mu\\,\\lsim$ 1 GeV which leads to large theoretical uncertainties in the quark star mass-radius relation. ",
        "introduction": "An understanding of the behavior of quantum chromodynamics (QCD) at high density is crucial to describing the physics of compact stars. This is due to the fact that the nuclear matter within these compact objects may be sufficiently dense to undergo a phase transition to a deconfined or ``quark-matter'' phase. Of particular interest is the possibility that stars which have a significant quark-matter component could have a dramatically different mass-radius relationship than normal neutron stars~\\cite{drake}. However, in order to make definitive statements about the mass-radius relationship for such objects one needs a reliable calculation of the equation of state of high-density QCD.  In recent years high-density QCD has received considerable attention due to the possible breaking of color gauge symmetry which gives rise to color superconductivity. As a consequence of this, QCD has a very complicated phase diagram which depends on the number and masses of the dynamical quarks~\\cite{frankie}. It would therefore seem that the effect of the superconducting phase on the equation of state would need to be taken into account; however, despite the fact that the gaps for color superconductivity are inherently non-perturbative, their effect on the equation of state for high-density QCD is expected to be small~\\cite{fraga}.  Therefore, it is possible to use equations of state that have been derived neglecting them.  The canonical choice for the equation of state for quark-matter has been to apply non-ideal bag models with various values for the bag constant~\\cite{bags}. However, others have applied quasiparticle models~\\cite{quasi} or directly applied the QCD weak-coupling expansion to first or second order in the strong coupling constant $\\alpha_s$~\\cite{fraga}.  The possibility of applying the weak-coupling expansion is an enticing option since this expansion can be unambiguously derived from first principles.  However, since the strong coupling constant is expected to be on the order of one in the phenomenologically relevant density range, the question of the convergence of the weak-coupling expansion becomes a very important one.  At high temperatures, for example, the weak-coupling expansion for the QCD equation of state  has been carried out to order $\\alpha_s^{5/2}$~\\cite{arnold1,kast,BN}.  Unfortunately, the finite-temperature weak-coupling expansion converges very slowly. In order to improve the convergence of the weak-coupling expansion calculational frameworks based on hard-thermal-loop resummation have been proposed~\\cite{kpp,EJM1,ejmp,spt,comp1,comp2}. These approaches attempt to describe finite-temperature QCD in terms of weakly-interacting massive quasiparticles by reorganizing the perturbative expansion around a state which includes hard-thermal-loop quasiparticles at lowest order.  Detailed studies using these techniques have shown that the convergence of the reorganized perturbative series appears to be much better than standard resummed perturbation theory\\cite{ejmp,spt}.  Motivated by the success of the hard-thermal-loop reorganization of perturbation theory we apply the equation of state for dense QCD at zero temperature obtained from hard-dense-loop perturbation theory (HDLpt) to quark stars.  The goal of doing this is to see if this technique can reduce the scale-dependence of the final results and improve the convergence of the successive approximations to the finite-density QCD equation of state. In this paper we calculate the equation of state of quark-matter to leading order in HDLpt and use the result to calculate the mass-radius relation for non-rotating quark stars.  We also make comparisons with the QCD weak-coupling expansion and discuss the large theoretical uncertainties in the weak-coupling and HDLpt results.  The paper is organized as follows. In section II, we discuss the weak-coupling expansion equation of state. In section III, we list the finite-temperature and density expressions for the leading order HTLpt/HDLpt free energy along with the HTL/HDL quark and gluon self-energies.  In section IV, we derive the leading order HDLpt equation of state at zero temperature and finite density. In section V, we use the resulting HDLpt equation of state to determine the mass-radius relationship of a non-rotating quark star. Finally, we summarize in section VI.  Necessary integrals are tabulated in the Appendix. ",
        "conclusions": "In this paper, we have calculated the free energy of cold dense quark matter to leading order in HDL perturbation theory (HDLpt). The predictions of HDLpt depend on a renormalization scale $\\Lambda$ that arises both from running of the coupling constant and from the renormalization of the additional ultraviolet divergences that are introduced by the HDLpt reorganization of perturbation theory. It is possible to separate these two effects by introducing two renormalization scales $\\Lambda_3$ and $\\Lambda_4$ as done in Ref.~\\cite{EJM1}, and these scales are associated with the soft scale and hard scale, respectively. For simplicity we have chosen not to distinguish between the two and have simply set $\\Lambda_3=\\Lambda_4$.  We then used the HDLpt equation of state as input to the TOV equations in order to determine the mass-radius relationship of quark stars.  We find that the large scale dependence of the HDLpt and weak-coupling expansion equations of state lead to large theoretical uncertainties in the quark star mass-radius relationship.  In the case of the weak-coupling expansion, the expansion only seems to be under control for $\\mu > 5$ GeV.  For HDLpt, the scale dependence is larger for all values of $\\mu$ and without a next-to-leading order HDLpt calculation it is not possible to draw conclusions about the convergence of the series. In addition, the choices $\\Lambda=4\\mu$ and $\\Lambda=2\\mu$ lead to rather unphysical predictions as can be seen in Fig.~\\ref{fig1}.  In order to eliminate these we were forced to further restrict the range of renormalization scales considered to $\\mu \\leq \\Lambda \\leq 1.6\\,\\mu$. The failure of both HDLpt and the weak-coupling expansion to reliably describe the finite-density QCD equation of state for $\\mu$ between 300 MeV and 1 GeV is troubling since this is the range which is important for determining the mass-radius relationship for a quark star.  As mentioned in Section \\ref{weaksec}, it possible that a computation of the order $\\alpha_s^3$ contribution to the finite-density QCD equation of state could remove some of the theoretical uncertainties resulting from the use of the weak-coupling expansion result.  Despite the fact that this would be a rather difficult task it seems that this calculation is required in order to draw more firm conclusions about the QCD equation of state. However, even if this calculation were available, the presence of a non-perturbative contribution from a color-superconducting phase of QCD in this range of quark chemical potential introduces additional theoretical uncertainties. Perturbative results extended down to this range of quark chemical potential give gaps on the order of $\\phi \\sim$ 30-100 MeV.  Since the gap gives a relative contribution of the order of $\\phi^2/\\mu^2$ this could translate into a relative modification of the equation of state between 1\\% and 10\\%.  A very challenging problem would be to calculate the next-to-leading order correction to the free energy in HDLpt.  If the next-to-leading order correction turns out to be small for relevant values of the chemical potential, the results obtained in the present work can be trusted.  However, it would seem more prudent to compute the order $\\alpha_s^3$ contribution in the weak-coupling expansion since the finite-density perturbation series does not seem to suffer from the same problems (oscillation and lack of convergence) as the finite-temperature perturbation series."
    },
    "0206/astro-ph0206189_arXiv.txt": {
        "abstract": "{ New diffraction-limited speckle images of the {\\rr} in the wavelength range 2.1--3.3\\,{\\mic} with angular resolutions of 44--68 mas \\citep{Tuthill_etal2002} and previous speckle images at 0.7--2.2\\,{\\mic} \\citep{Osterbart_etal1997,Men'shchikov_etal1998} revealed well-resolved bright bipolar outflow lobes and long {\\sf X}-shaped spikes originating deep inside the outflow cavities. This set of high-resolution images stimulated us to reanalyze all infrared observations of the {\\rr} using our two-dimensional radiative transfer code. The high-resolution images imply a geometrically and optically thick torus-like density distribution with bipolar conical cavities and are inconsistent with the flat disk geometry frequently used to visualize bipolar nebulae. The new detailed modeling, together with estimates of the interstellar extinction in the direction of the {\\rr} enabled us to more accurately determine one of the key parameters, the distance $D \\approx$ 710\\,pc with model uncertainties of 70\\,pc, which is twice as far as the commonly used estimate of 330\\,pc. The central binary is surrounded by a compact, massive ($M \\approx$ 1.2\\,{\\Msun}), very dense dusty torus with hydrogen densities reaching $n_{\\rm H} \\approx 2.5 \\times 10^{12}$\\,cm$^{-3}$ (dust-to-gas mass ratio {\\dustgas} $\\approx 0.01$). The model implies that most of the dust mass in the dense torus is in very large particles and, on scales of more than an arcsecond, the polar outflow regions are denser than the surrounding medium. The bright component of the spectroscopic binary {\\hd} is a post-AGB star with mass {\\Mstar} $\\approx 0.57$\\,{\\Msun}, luminosity {\\Lstar} $\\approx 6000$\\,{\\Lsun}, and effective temperature {\\Tstar} $\\approx 7750$\\,K. Based on the orbital elements of the binary, we identify its invisible component with a helium white dwarf with {\\Mwd} $\\approx$ 0.35\\,{\\Msun}, {\\Lwd} $\\sim$ 100\\,{\\Lsun}, and {\\Twd} $\\sim 6 \\times 10^{4}$\\,K. The hot white dwarf ionizes the low-density bipolar outflow cavities inside the dense torus, producing a small H\\,II region observed at radio wavelengths. We propose an evolutionary scenario for the formation of the {\\rr} nebula, in which the binary initially had 2.3 and 1.9\\,{\\Msun} components at a separation of $\\sim$ 130\\,{\\Rsun}. The nebula was formed in the ejection of a common envelope after Roche lobe overflow by the present post-AGB star. \\keywords { radiative transfer -- circumstellar matter -- stars: individual: {\\rr} -- stars: mass-loss -- stars: AGB and post-AGB -- infrared: stars } } ",
        "introduction": "\\label{Introduction}  The {\\rr} is a spectacular bipolar reflection nebula around an evolved close binary star (also known as {\\hd}, \\object{AFGL\\,915}, \\object{IRAS\\,06176--1036}). The object has been extensively studied for more than two decades \\citep[see, e.g., references in][]{Waters_etal1998,Men'shchikov_etal1998}. Recent diffraction-limited speckle images of the object with 62--76\\,mas resolution in the optical \\citep[0.6--0.8\\,{\\mic},][ hereafter Paper~I]{Osterbart_etal1997} and near infrared \\citep[0.7--2.2\\,{\\mic},][ hereafter Paper~II]{Men'shchikov_etal1998} displayed a compact, highly symmetric bipolar nebula with pronounced {\\sf X}-shaped spikes, implying a toroidal distribution of the circumstellar material. No direct light from the completely obscured central binary could be seen.  New diffraction-limited images of the {\\rr} in the near-IR (2.1--3.3\\,{\\mic}) with unprecedented angular reso\\-lutions of 46--68\\,mas were recently presented by \\cite{Tuthill_etal2002} (hereafter Paper~III). The images were reconstructed from the Keck telescope speckle data using the bispectrum speckle interferometry method. The highest-resolution images clearly show two bright lobes above and below the midplane of an inclined torus or geometrically very thick disk. {\\sf X}-shaped spikes along the surfaces of the conical outflow cavities contribute to the intensity distribution of the two bright lobes, making them appear widened and even double-peaked. A striking feature of the {\\rr} bipolar nebula is its self-similar appearance on scales from 80\\,mas to 1{\\arcmin} and from the red light to at least 10\\,{\\mic}, implying that large grains of at least several microns in size dominate scattering.  \\begin{figure*} \\hspace{0.075\\hsize} \\resizebox{0.85\\hsize}{!} { \\hspace{-5mm} \\includegraphics{h3541f01.eps} \\hspace{10mm} \\includegraphics{h3541f02.eps} } \\caption { Geometry of the circumstellar environment of the close binary {\\hd} ({\\em left panel}) and three-dimensional representation of the massive circumbinary torus of the {\\rr} ({\\em right panel}) as it appears in the near-IR images of \\citetalias{Osterbart_etal1997}, \\citetalias{Men'shchikov_etal1998}, and \\citetalias{Tuthill_etal2002} in projection onto the sky plane. Schematically shown are four regions of the model -- the innermost dense torus with bipolar cavities (100\\,AU radius; dark color), the less dense envelope with a ${\\rho} \\propto r^{-1.5}$ density profile (400\\,AU radius; medium color), the bipolar outflow cavities (lighter color), and the outer extended envelope with a steep ${\\rho} \\propto r^{-4}$ density gradient (4$ \\times 10^{4}$\\,AU; the lightest color). The geometry is defined by the opening angle of the cavities, $\\omega = \\pi - \\psi \\approx$ 50{\\degr} ($\\psi \\approx$ 130{\\degr}) and the viewing angle, $\\theta_{\\rm v} \\approx$ 11{\\degr}, between the equatorial plane and the line of sight. } \\label{Geometries} \\end{figure*}  Only a few two-dimensional (2D) radiative transfer calculations of the bipolar envelope of the {\\rr} have been published to date. \\citet{Yusef-Zadeh_etal1984} first simulated the well-known optical images of the nebula with {\\sf X}-shaped spikes using a Monte-Carlo scattering method. Varying the density distribution and scattering properties of dust grains, they found that a torus-like configuration with a $\\rho \\propto r^{-2}$ radial density profile and with biconical cavities having a full opening angle of 70{\\degr} is able to reproduce the shape of the nebula. These calculations gave support to the previously suggested idea \\citep{Cohen_etal1975,Morris1981,Perkins_etal1981} that a quasi-spherical envelope with bipolar cavities can reproduce the shape of the {\\rr}.  Extending the previous modeling, \\citet{Lopez_etal1997} applied a Monte-Carlo technique in which they were able to not only simulate scattering of the stellar radiation at a selected wavelength, but also calculate the radiative equilibrium temperature and emission of dust. With multi-wavelength radiative transfer computations, they aimed to constrain the model by comparing it with the observed spectral energy distribution (SED) and a deconvolved 0{\\farcs}2 resolution intensity map at 2.2\\,{\\mic} \\citep{Cruzalebes_etal1996}. This work represented the first step in the direction of a more realistic modeling of the {\\rr} capable of explaining a larger set of observational data and to reconstruct reliable properties of the object.  In the previous 2D modeling presented in \\citetalias{Men'shchikov_etal1998}, we applied a frequency-dependent ray tracing method \\citep{Men'shchikovHenning1997} to construct a detailed model of the {\\rr} consistent with a much larger number of observational constraints. For the first time, the model reproduced reasonably well the entire SED of the {\\rr} from the ultraviolet to centimeter wavelengths and the highest-resolution (76 mas) speckle-interferometry images at 0.656\\,{\\mic}, 0.8\\,{\\mic}, 1.65\\,{\\mic}, and 2.2\\,{\\mic}. The extensive modeling allowed us to derive the geometry of a compact circumbinary torus-like structure, such as the opening angle $\\omega = 70${\\degr} of the bipolar cavities and the inclination angle $\\theta_{\\rm v} = 7${\\degr} of the symmetry axis. For an assumed distance of 330 pc, the model reconstructed physical parameters of the object, such as the total luminosity $L_{\\star} \\approx 3000$ {\\Lsun}, the radius $R \\sim 30$ AU and the mass $M \\approx 0.25$ {\\Msun} of the opaque ($A_V \\approx 30$) torus, the density distribution $\\rho \\propto r^{-2}$ for $r < 16$ AU and $\\rho \\propto r^{-4}$ for $r > 16$ AU, and very large sizes of dust particles ($a \\sim 0.2$ cm).  Although the model describes reasonably well the large number of constraints, our new Keck telescope images of the {\\rr} with unprecedented resolutions \\citepalias[44--68 mas,][]{Tuthill_etal2002} have shown that the model is not consistent with the longest-wavelength 3.1\\,{\\mic} and 3.3\\,{\\mic} images. As we have already demonstrated in \\citetalias{Tuthill_etal2002}, the model predicts a single elongated peak at this wavelength, whereas the Keck image clearly displays two lobes divided by a dark lane, very similar to the shorter-wavelength images. In the model, there is too much emission from the hot grains close to the inner boundary of the torus, which implies insufficient optical depths. This finding stimulated us to recompute the model taking into account the new constraints in addition to the old data.  This paper presents a new detailed study of the {\\rr} based on our previous model \\citepalias{Men'shchikov_etal1998} and on the new constraints provided by the Keck telescope images \\citepalias{Tuthill_etal2002}. In Sect.~\\ref{RadTraModel} we describe our assumptions and radiative transfer model of the dusty circumbinary torus. In Sect.~\\ref{ModelResults} we present the model results and compare them with available observational data. In Sect.~\\ref{Discussion} we discuss the parameters of the {\\rr} and evolution of its close binary. In Sect.~\\ref{Conclusions} we summarize conclusions of this work. ",
        "conclusions": "\\label{Conclusions}  Recent diffraction-limited near-IR speckle images of the {\\rr} in the wavelength range 2.1--3.3\\,{\\mic} with angular resolutions of 44--68 mas \\citepalias{Tuthill_etal2002} as well as the previous optical and near-IR speckle images at 0.7--2.2\\,{\\mic} \\citepalias{Osterbart_etal1997,Men'shchikov_etal1998} revealed a geometrically thick circumbinary torus with bipolar outflow cones and {\\sf X}-shaped spikes originating deep inside the cavities. This multiwavelength set of high-resolution images enabled us to reanalyze most of the existing observations using our two-dimensional radiative transfer code. Results of this study are summarized below.  {\\sc Distance.} An important by-product of the modeling of the {\\rr} is a new determination of its distance $D \\approx 710$\\,pc (with model uncertainties of about 10\\,{\\%}), which is twice as large as the commonly adopted value of 330\\,pc (Sect.~\\ref{Distance}). The new distance, based on the account for interstellar extinction in our model (Sects.~\\ref{Distance}, \\ref{SpEnDi}), is consistent with high luminosities expected from the stellar evolution theory. Based on this distance, we reconstructed physical parameters of the binary and its circumbinary torus (Table~\\ref{ModelParams}).  {\\sc Close binary.} In our model, the observed component of the spectroscopic binary is a luminous, low-mass post-AGB star with {\\Mstar} $\\approx 0.57$\\,{\\Msun}, {\\Tstar} $\\approx 7750$\\,K, and {\\Lstar} $\\approx 6050$\\,{\\Lsun}. It is a product of the mass loss by the secondary component in a close binary system. We identified the now invisible descendant of the primary component with a relatively hot white dwarf with {\\Mwd} $\\approx$ 0.35\\,{\\Msun}, {\\Twd} $\\approx 6 \\times 10^{4}$\\,K, and {\\Lwd} $\\approx$ 100\\,{\\Lsun}. The presence and parameters of the compact degenerate star in the close binary {\\hd} have been deduced (Sect.~\\ref{Binary}) from (1) the spectroscopic mass function $f(M)$ = 0.049\\,{\\Msun}, (2) presence of a compact H\\,II region ionized by a hot source of radiation, (3) upper limit of the contribution of the hot object to the continuum in the far UV.  {\\sc Circumbinary structure.} The intensity distribution of the high-resolution images is definitely inconsistent with the flat disk geometry frequently used to visualize bipolar nebulae (Sect.~\\ref{Bipolar}). A geometrically very thick density distribution of a compact, dense torus with biconical outflow cavities (Fig.~\\ref{Geometries}) is best suitable for reproducing the observed images. The opening angle of the cavities is 50{\\degr} and the observer's viewing angle is 11{\\degr} below the midplane.  {\\sc Density distribution.} Although the nebula extends to at least $R_2 \\approx 4 \\times 10^{4}$ AU from the central binary, most of its mass of $M \\approx 1.2$\\,{\\Msun} is contained in the extremely dense torus having an outer radius of 100\\,AU (Sect.~\\ref{DensTemp}). The inner {\\em dust} boundary of the torus is located at a distance $R_1 \\approx 14$\\,AU from the center, where gas densities reach values of $\\rho \\approx 4.2 \\times 10^{-12}$ g\\,cm$^{-3}$ ($n_{\\rm H} \\approx 2.5 \\times 10^{12}$\\,cm$^{-3}$). The density of the outflow cavities is many orders of magnitude lower than the torus density in the region dominated by the dense torus (Fig.~\\ref{DenTem}), whereas in the outer regions ($r \\ge$ 800 AU) the outflow cones are denser than the rest of the toroidal envelope.  {\\sc Dust properties.} Our model of dust in the circumbinary torus of the {\\rr} has two distinct components whose chemical composition is not constrained by observations. Amorphous carbon grains with radii $a$ in the range of 0.005--600\\,{\\mic} and typical interstellar size distribution ${\\rm d}n/{\\rm d}a \\propto a^{-3.5}$ exist mainly in the outer regions of the toroidal envelope ($r > 100$\\,AU). Most of the dust mass is contained in very large ($a$ = 0.2\\,cm) particles of the massive, dense torus, which produce an almost gray optical depth of $\\tau \\approx 47$. Although there exist more dust components somewhere in the nebula, including crystalline silicates and PAHs, their mass must be small compared to the mass of the very large particles, which dominate the observed appearance of the {\\rr}.  {\\sc Evolution of the binary.} Based on our estimates of the masses of the {\\rr} binary components, their separation and luminosities, we suggest an evolutionary scenario for the formation of the nebula, in which components had initial masses of about 2.3 and 1.9\\,{\\Msun} and a separation of $\\sim 130\\,${\\Rsun}. The scenario associates the formation of the {\\rr} nebula with the ejection of the common envelope upon the Roche lobe overflow by the present post-AGB star.   \\begin{acknowledgement} We are grateful to Viktor Malanushenko for the assistance in computing the orbital parameters, to Anatoly Miroshnichenko for his help in the derivation of the observational estimate of the interstellar reddening, and to Jarrod Hurley for providing the SSE package. LRY was supported by the Russian ``Astronomy and Space Research'' program. This research has made use of the SIMBAD database operated at CDS, Strasbourg, France, and of the data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center, funded by the National Aeronautic and Space Administration and the National Science Foundation. We thank the referee, Michael Barlow, for his very useful comments. \\end{acknowledgement}"
    },
    "0206/astro-ph0206376_arXiv.txt": {
        "abstract": "SAX J1808.4--3658 is a unique source being the first Low Mass X--ray Binary showing coherent pulsations at a spin period comparable to that of millisecond radio pulsars. Here we present an XMM-Newton observation of SAX J1808.4--3658 in quiescence, the first which assessed its quiescent luminosity and spectrum with good signal to noise. XMM-Newton did not reveal other sources in the vicinity of SAX J1808.4--3658 likely indicating that the source was also detected by previous BeppoSAX and ASCA observations, even if with large positional and flux uncertainties. We derive a 0.5--10 keV unabsorbed luminosity of $L_X=5\\times10^{31}\\ergs$, a relatively low value compared with other neutron star soft X--ray transient sources. At variance with other soft X--ray transients, the quiescent spectrum of SAX J1808.4--3658 was dominated by a hard ($\\Gamma\\sim 1.5$) power law with only a minor contribution ($\\lsim 10\\%$) from a soft black body component. If the power law originates in the shock between the wind of a turned-on radio pulsar and matter outflowing from the companion, then a spin-down to X--ray luminosity conversion efficiency of $\\eta\\sim 10^{-3}$ is derived; this is in line with the value estimated from the eclipsing radio pulsar PSR J1740--5340. Within the deep crustal heating model, the faintness of the blackbody-like component indicates that SAX J1808.4--3658 likely hosts a massive neutron star ($M\\gsim1.7\\msole$). ",
        "introduction": "Neutron star Soft X--ray Transients (SXRTs) when in outburst closely resemble persistent Low Mass X--ray Binaries (LMXRBs). In the last few years it has become clear that SXRT sources form a rather inhomogeneous class (see Campana et al. 1998a for a review) comprising sources with well defined outbursts as well as sources with long on/off activity periods. Moreover, sources displaying bright outbursts with peak X--ray luminosities $L_X\\sim 10^{37}-10^{38}$~erg~s$^{-1}$ appear to be different from sources showing only faint outbursts reaching $L_X\\sim 10^{36}-10^{37}$~erg~s$^{-1}$, especially in the Galactic center region (Heise et al. 1998; King 2000; in't Zand 2001a).  A major leap forward came with the discovery of SAX J1808.4--3658, a bursting SXRT reaching a maximum luminosity of $\\sim 2\\times 10^{36}\\ergs$ (for a distance of 2.5 kpc; in't Zand al. 2001b). In April 1998 the source resumed activity and RXTE observations revealed coherent $\\sim 401$~Hz pulsations, the first detected in the persistent emission of a neutron star LMXRB. These testify to the presence of magnetic polar cap accretion onto a fast rotating magnetic neutron star (Wijnands \\& van der Klis 1998; Chakrabarty \\& Morgan 1998). The inferred magnetic field strength of SAX~J1808.4--3658 is in the $10^8-10^9$~G range (see Psaltis \\& Chakrabarty 1999), providing convincing evidence for the long suspected LMXRB-millisecond pulsar connection.  SXRTs spend much of their time in quiescence. The origin of the quiescent X--ray emission is still uncertain. In the last few years several sources have been studied in detail; the picture that emerged is that the quiescent spectrum consists of a soft component plus a high energy excess. The former component is often fit with a blackbody spectrum with an equivalent radius of 1--2~km and temperatures in the 0.1--0.3~keV range or with a neutron star atmosphere model with an equivalent radius consistent with the entire neutron star surface and slightly smaller temperatures (Brown et al. 1998; Rutledge et al. 2000). The high energy component is well represented by a power law with photon index $\\Gamma\\sim 1-2$ (Asai et al. 1996, 1998; Campana et al. 1998b, 2000). In all sources observed so far the quiescent luminosity ranges between $10^{32}-10^{33}\\ergs$, indicating a clear difference with black hole transients in quiescence that have a lower X--ray luminosity (Garcia et al. 2001; Campana \\& Stella 2000). Flux and spectral variability have been reported in Aql X-1 and KS 1731--260 during quiescence (Campana et al. 1997; Rutledge et al. 2002; Wijnands et al. 2002a). This poses severe limitations on the emission mechanisms responsible for the quiescent luminosity.  Among neutron star SXRTs, SAX J1808.4--3658 stands out for having, while in outburst, a magnetosphere and, of course, an accurately measured spin period. These characteristics make the source very well suited for testing the predictions of models for the quiescent emission in which the presence of a sizeable magnetic field plays a crucial role (Stella et al. 1994). SAX J1808.4--3658 was detected in quiescence (Stella et al. 2000; Dotani et al. 2001; Wijnands et al. 2001) though with large uncertainties. Here we report on an XMM-Newton observation of SAX~J1808.4--3658 in quiescence, the first to detect the source with a good signal to noise ratio. ",
        "conclusions": "The high throughput and good angular resolution of XMM-Newton provided the first firm determination of the quiescent luminosity of SAX J1808.4--3658. This is lower by a factor of $\\sim 2$ than previous best fit estimates. However, once the large uncertainties in the spectral parameters are taken into account, a fairly constant luminosity is inferred since March 1999 at a level of $5\\times 10^{31}\\ergs$ (at a distance of 2.5 kpc; note that the reduced $\\chi^2$ of a fit with a constant is 1.8, $15\\%$ null hypothesis variability, see Fig. 3). We conclude that the source variability issue raised by Dotani et al. (2000) is questionable. The source luminosity is a factor of two lower than that usually measured in SXRTs in quiescence; this ranges between $10^{32}-10^{33}\\ergs$ (e.g. Campana et al. 1998a; Campana \\& Stella 2000; Garcia et al. 2001). For a distance of 4 kpc, however, it would be fully consistent with what usually observed in other SXRTs.  More interestingly, if we consider the spectral fit with a power law component plus a black body, the soft X--ray component comprises only a small part ($\\lsim 10\\%$) of the total luminosity, whereas in the great majority of SXRTs it accounts for about half of the luminosity in the 0.5--10 keV energy band (e.g. Campana et al. 1998b, 2000; see however Rutledge et al. 2002). The soft component is usually ascribed to the cooling of the neutron star surface powered by the deep nuclear heating that the neutron star receives during each outburst (Brown et al. 1998; Rutledge et al. 2000; Colpi et al. 2001; see also Campana et al. 1998a). RXTE All Sky Monitor (ASM) provides us with a continuous monitoring of the high level activity of SAX J1808.4--3658. From these data we can extrapolate a mean mass transfer rate of $\\sim 5\\times 10^{-12}\\msole$ yr$^{-1}$, in line with previous estimates (Bildsten \\& Chakrabarty 2001). This mean mass inflow rate translates into a soft quiescent luminosity of $10^{32}\\ergs$ within the 0.5--10 keV energy band from deep crustal heating (Brown et al. 1998; Colpi et al. 2001). This is a factor of $\\sim 10$ higher than observed. Given the fact that predictions from deep nuclear heating are quite robust, one is led to conclude that an additional source of cooling is present. A simple and well known solution is when the direct Urca process is allowed in the neutron star core and, in turn, neutrino cooling does affect the neutron star thermal evolution (Colpi et al. 2001). This can occur only for massive neutron stars with masses higher than $1.7-1.8\\msole$. If this interpretation were correct the neutron star of SAX J1808.4--3658 have to be fairly massive, in agreement with accretion spin-up scenarios.  The main contribution to the quiescent luminosity of SAX J1808.4--3658 appears to derive from the power law component. A pure propeller contribution is ruled out since this mechanism should stop operating at a luminosity of about $10^{33}\\ergs$, with the turing on of a radio pulsar (e.g. Campana et al. 1998a, 1998b). One interpretation for this power law relies on the emission at the shock front between the relativistic wind of a radio pulsar and matter outflowing from the companion star (see the discussion in Stella et al. 2000). Being the first SXRT for which the presence of a sizable magnetic field is unambiguously established, SAX J1808.4--3658 can be used to infer the efficiency $\\eta$ with which spin-down power is converted into 0.5--10 keV luminosity. We derive $\\eta\\sim 5\\times 10^{-3}\\,B_8^2$, where $B_8=B/10^8\\, {\\rm G}$ is neutron star magnetic field (as derived from the accretion luminosity at the propeller onset, Gilfanov et al. 1998; see also Psaltis \\& Chakrabarty 1999).  Recently, an ideal laboratory for studying the shock emission has been discovered: this is PSR J1740--5340, a 3.7 ms radio pulsar in a 32.5 hr orbit around a Roche lobe filling main sequence companion, which display partial and total eclipsed over a wide range of orbital phases (D'Amico et al. 2001; Ferraro et al. 2001; Burderi, D'Antona \\& Burgay 2002). The X--ray luminosity of  PSR J1740--5340 measured by Chandra ($8\\times 10^{30}\\ergs$, 0.5--2.5 keV range, unabsorbed; Grindlay et al. 2001) implies $\\eta\\sim 10^{-4}$. Taking this system as a prototype for modelling the shock emission power (which is likely less efficient since the radio pulsar is not  completely engulfed) and scaling the efficiency as the square of the orbital separation, one infers $\\eta\\sim 3\\times 10^{-3}$, well consistent with the value above, given the uncertainties involved."
    },
    "0206/gr-qc0206041_arXiv.txt": {
        "abstract": "Gravitational wave emission from the gravitational collapse of massive stars has been studied for more than three decades.  Current state of the art numerical investigations of collapse include those that use progenitors with more realistic angular momentum profiles, properly treat microphysics issues, account for general relativity, and examine non--axisymmetric effects in three dimensions.  Such simulations predict that gravitational waves from various phenomena associated with gravitational collapse could be detectable with ground--based and space--based interferometric observatories. ",
        "introduction": "\\label{section:introduction}  The field of gravitational wave (GW) astronomy will soon become a reality. The first generation of ground--based interferometric detectors (LIGO~\\cite{ligo}, VIRGO~\\cite{virgo}, GEO 600~\\cite{geo600}, TAMA 300~\\cite{tama300}) are beginning their search for GWS.  Towards the end of this decade, two of these detectors (LIGO, VIRGO) will undergo upgrades that should allow them to reach sensitivities necessary to regularly detect emission from astrophysical sources.  A space--based interferometric detector, LISA~\\cite{lisa}, could be launched in the early part of the next decade. One important class of sources for these observatories is stellar gravitational collapse.  This class includes the accretion induced collapse (AIC) of white dwarf binary components and the core collapse of massive stars ($M>8 M_{\\odot} $), very massive Population III stars ($M$=$100$-$500 M_{\\odot} $), and supermassive stars (SMSs, $M>10^6 M_{\\odot}$).  Some of these collapses result in explosions (Type II, Ib/c supernovae and hypernovae) and all leave behind neutron star or black hole remnants.  Strong GWs can be emitted during a gravitational collapse/explosion and, following the collapse, by the resulting compact remnant~\\cite{thorne-94, mueller-97a, mueller-98, finn-99, schutz-99, fryer-02, fryer-02b, hughes-02}.  GW emission during the collapse itself may result if the collapse or explosion involves aspherical bulk mass motion or convection.  Rotational or fragmentation instabilities encountered by the collapsing star will also produces GWs.  Neutron star remnants of collapse may emit GWs due to the growth of rotational or r-mode instabilities.  Black hole remnants will also be sources of GWs if they experience accretion induced ringing.  All of these phenomena have the potential of being detected by gravitational wave observatories because they involve the rapid change of dense matter distributions.  Observation of gravitational collapse by gravitational wave detectors will provide unique information, complementary to that derived from electromagnetic and neutrino detectors. Gravitational radiation arises from the coherent superposition of mass motion.  Whereas, electromagnetic emission is produced by the incoherent superposition of radiation from electrons, atoms, and molecules.  Thus GWs carry different kinds of information than other types of radiation.  Furthermore, electromagnetic radiation interacts strongly with matter and thus only gives a view of the collapse from lower density regions near the surface of the star and is weakened by absorption as it travels to the detector.  Neutrinos can escape from much further within the collapsing star, but even they are scattered by the highest density regions in the core. By contrast, gravitational waves can propagate from the innermost parts of the stellar core to detectors without attenuation by intervening matter.  The characteristics of the GW emission from gravitational collapse have been the subject of much study.  Core collapse supernovae, in particular, have been investigated as sources of gravitational radiation for more than three decades (see, e.g.,~\\cite{ruffini-71, thuan-74, saenz-78, detweiler-81, nakamura-81, mueller-82, stark-85, finn-90, monchmeyer-91, zwerger-97, rampp-98, fryer-02, fryer-02b}). However, during this time research has produced estimates of GW strength that vary over orders of magnitude.  This is due to the complex nature of core collapse. Important theoretical and numerical issues include \\begin{itemize} \\item construction of accurate progenitor models, including realistic angular momentum distributions, \\item proper treatment of microphysics, including the use of realistic equations of state and neutrino transport, \\item simulation in three-dimensions to study non-axisymmetric effects, \\item inclusion of general relativistic effects, \\item inclusion of magnetic field effects, and \\item study of the effect of an envelope on core behavior. \\end{itemize} To date, collapse simulations generally include state of the art treatments of only one or two of the above physics issues (often because of numerical constraints).  For example, those studies that include advanced microphysics have often been run with Newtonian gravity (and approximate evaluation of the GW emission; see section~\\ref{section:aic-num}).  A 3D, general relativistic collapse simulation, which includes all significant physics effects, is not feasible at present.  However, good progress has been made on the majority of the issues listed above; the more recent work will be reviewed in some detail here.  The remainder of this article is structured as follows. Each category of gravitational collapse will be discussed in a separate section (AIC in section~\\ref{section:aic}, collapse of massive stars in section~\\ref{section:sne}, collapse of Population III stars in section~\\ref{section:pop3}, and collapse of SMSs in section~\\ref{section:sms}).  Each of these sections (\\ref{section:aic}, \\ref{section:sne}, \\ref{section:pop3}, \\ref{section:sms}) is divided into subsection topics: Collapse Scenario, Formation Rate, GW Emission Mechanisms, and Numerical Predictions of GW Emission.  In the subsections on numerical predictions, the detectability of the GW emission from various phenomena associated with collapse is examined.  In particular, the predicted characteristics of GW emission are compared to the sensitivities of LIGO (for sources with frequencies of $1$ to $10^4\\,{\\rm Hz}$) and LISA (for sources with lower frequencies in the range of $10^{-4}$ to $1\\,{\\rm Hz}$).  \\newpage ",
        "conclusions": "\\label{section:summary}  It is hoped that as gravitational collapse simulations become more sophisticated, the historically widely varying estimates of the magnitude of GW emission from collapse may start to converge.  Steady progress in this field has been made in the last decade.  Some researchers have begun to use progenitor models produced with stellar evolution codes, which thus have more realistic angular momentum profiles, as starting points for collapse simulations~\\cite{fryer-02}.  This reduces the need for collapse studies that include large surveys of the angular momentum parameter space.  Other progress made in the numerical study of collapse includes the use of realistic equations of state~\\cite{fryer-02}, advanced neutrino transport and interaction schemes~\\cite{janka-02c, lieben-02, rampp-02, thompson-02}, and the performance of 3D Newtonian~\\cite{rampp-98, brown-01, fryer-02a} and improved 2D general relativistic simulations~\\cite{shibata-00ptp}.  There is still much work to be done toward the goal of self--consistent, 3D general relativistic collapse simulations. Accurate progenitor modelling and collapse simulations must include the effects of magnetic fields, as they can significantly alter the amount of angular momentum and differential rotation present in collapsing stars.  Many of the more advanced studies, which include proper microphysics treatment and/or general relativistic effects, have been limited to axisymmetry.  Full 3D simulations are necessary to compute the characteristics of the GW emission from non--axisymmetric collapse phenomena.  Furthermore, simulations that follow both the collapse and the evolution of the collapsed remnant are necessary to consistently predict GW emission.  One benefit of long duration simulations is that they will facilitate the investigation of the effects of the envelope on any instabilities that develop in the collapsing core or remnant.  Of course, lengthy 3D simulations are computationally intensive.  This burden may be reduced by the use of advanced numerical techniques, including adaptive mesh refinement and parallel algorithms.  The current numerical simulations of gravitational collapse indicate that interferometric observatories could detect GWs emitted by some collapse phenomena.  LIGO-I may be able to detect GWs from secular bar--mode instabilities in core--collapse SNe~\\cite{lai-01} and magnetized tori surrounding black hole collapse remnants~\\cite{putten-02}.  LIGO-II could observe GWs from dynamical bar--mode instabilities in AIC~\\cite{liu-02} and core--collapse SNe~\\cite{fryer-02}, and possibly from the fragmentation of very massive SNe cores that merge to form BHs~\\cite{davies-02}.  LISA should be able to detect the collapse (and any bar--mode instabilities that develop during the collapse) of SMSs~\\cite{baumgarte-99} and the ringdown of black hole remnants of collapsed Population III stars~\\cite{fryer-02} and SMSs~\\cite{saijo-02}.  These observations will provide unique information about gravitational collapse and its associated progenitors and remnants.  \\clearpage \\newpage"
    },
    "0206/astro-ph0206006_arXiv.txt": {
        "abstract": "In standard CDM halo models, the time delay of a gravitational lens is determined by the cold baryon mass fraction, $f_b=\\Omega_{b,cold}/\\Omega_0$, of the visible galaxy relative to the overall halo.  The observed time delays in PG1115+080, SBS1520+530, B1600+434 and HE2149--2745 give Hubble constants consistent with the HST Key Project value of $H_0=72 \\pm 8$~km/s~Mpc only if $f_b \\gtorder 0.2$ (1-sided 68\\% confidence), which is larger than the upper bound of $f_{b,max}=\\Omega_b/\\Omega_0=0.15\\pm0.05$ estimated from the CMB.  If all available baryons cool and $f_b=f_{b,max}$ then the time delays imply $H_0=65 \\pm 6$~km/s~Mpc (95\\% confidence). If local inventories of cold baryons, $f_b\\simeq 0.013/h_{70}$, are correct, then $H_0=52\\pm 6$~km/s~Mpc and the halo parameters closely match isothermal mass models.  Isothermal models are also consistent with strong and weak lens studies, stellar dynamics and X-ray observations on these scales, while significantly more centrally concentrated models are not. There is a a conflict between gravitational lens time delays, the local distance scale and standard CDM halo models. ",
        "introduction": "Kochanek~(\\cite{Kochanek02a}) found that it was difficult to reconcile the time delays measured for 5 simple, well-observed gravitational lenses with the local distance scale given our expectation that galaxies have massive, extended dark matter halos.  If the lens galaxies had constant mass-to-light ($M/L$) ratios we found $H_0=71\\pm6$~km/s~Mpc, which is consistent with the local estimate of $H_0=72\\pm8$~km/s~Mpc by the HST Key Project (Freedman et al.~\\cite{Freedman01}).  However, if the lenses had isothermal mass distributions (flat rotation curves), we found $H_0=48_{-4}^{+7}$~km/s~Mpc, which is grossly inconsistent with the HST Key Project.  While the time delay lenses cannot distinguish between these two limiting mass distributions, models of other lenses (e.g. Munoz, Kochanek \\& Keeton~\\cite{Munoz01}), stellar dynamical measurements (e.g. Rix et al.~\\cite{Rix97}, Romanowsky \\& Kochanek~\\cite{Romanowsky99}, Gerhard et al.~\\cite{Gerhard01}, Treu \\& Koopmans~\\cite{Treu02}), weak lensing (e.g. Guzik \\& Seljak~\\cite{Guzik02}) and X-ray (e.g. Fabbiano~\\cite{Fabbiano89}, Lowenstein \\& White~\\cite{Lowenstein99}) measurements all suggest that the isothermal mass distributions are correct.  In this study we will show that standard cold dark matter (CDM) halo models closely resemble the isothermal models on these scales, which implies there is a conflict between the local distance scale, gravitational lens time delays and CDM halo models.  While the Kochanek~(\\cite{Kochanek02a}) results provided evidence for a real conflict given the considerable observational evidence that lens galaxies must have extended, massive dark matter halos, the link to a problem with CDM halo models was qualitative because the study lacked a quantitative, theoretical prediction for the time delays expected from CDM halos.  One barrier to making such predictions was that we lacked a clear understanding of which features of lens mass distributions control time delays.  While global degeneracies due to the addition of constant mass density sheets (e.g. Falco, Gorenstein \\& Shapiro~\\cite{Falco85}, Gorenstein, Falco \\& Shapiro~\\cite{Gorenstein88}, Saha~\\cite{Saha00}) and a correlation between more compact mass distributions and longer time delays (e.g. Schechter~\\cite{Schechter00}, Witt, Mao \\& Keeton~\\cite{Witt00}) were well known, it was unclear which properties of a halo model had to be accurately computed in order to make robust predictions. Kochanek~(\\cite{Kochanek02b}) combined analytic results with comparisons to the numerical models by Kochanek~(\\cite{Kochanek02a}) to show that the surface density in the annulus between the images used to measure the delay was the most important physical property of the lens galaxy for determining the time delay. The interior mass is implicit in the astrometry of the images and the lens galaxy, and the angular structure is either unimportant or strongly constrained by the astrometry. As a result, the Hubble constant expected for a simple lens is related to the surface density by $H_0 = A(1-\\kbar) + B(\\eta-1)\\kbar$, where $\\kbar$ is the average surface density in the annulus between the images (in units of the critical density), with a modest correction $|B| \\ltorder A/10$ due to the logarithmic slope $\\eta$ of the surface density distribution within the annulus ($\\kappa \\propto R^{1-\\eta}$). The coefficients $A$ and $B$ depend only on the image positions and the measured time delay.  These simple semi-analytic scaling laws reproduce full numerical models to accuracies of better than 5\\%.  We can now calculate the expected properties of gravitational lens time delays for CDM halo models.  In \\S2 we outline our model for the halos, which are based on the CDM lens models from Keeton~(\\cite{Keeton01}). The models consist of a Hernquist~(\\cite{Hernquist90}) model for the luminous early-type lens galaxy embedded in an NFW (Navarro, Frenk \\& White~\\cite{Navarro96}) halo normalized using the parameter estimates of Bullock et al.~(\\cite{Bullock01}). We considered both unmodified NFW halos and adiabatically compressed (Blumenthal et al.~\\cite{Blumenthal86}) halos.  We summarize the mathematical details of the model in \\S2.   In \\S3 we apply it to the four simple time delay lenses PG1115+080, SBS1520+530, B1600+434 and HE2149--2745, to show that the values of $\\kbar$ and $\\eta$ that determine the Hubble constant given the measured time delays are in turn determined by a single parameter, the cold baryonic mass fraction, $f_b = \\Omega_{b,cold}/\\Omega_0$, of the luminous galaxy compared to the halo.  Since the baryon fraction is bounded by local estimates from observed baryonic populations and the global baryon fraction estimated either in clusters or from the CMB, we can set firm bounds for the range of $H_0$ consistent with CDM halo models. As we discuss in \\S4, this leads to a new element of the so-called ``dark matter crisis'' (e.g. Moore~\\cite{Moore01}), because the CDM halo models combined with the measured time delays require lower Hubble constants than are consistent with the Key Project estimates based on the local distance scale.  An Appendix briefly discusses the effects of tidal truncation on lens galaxy halos.  \\def\\bfx{{\\bf x}} \\def\\bfu{{\\bf u}} \\def\\grad{{\\bf \\nabla}} \\def\\ka{\\kappa_1} \\def\\kb{\\kappa_2} ",
        "conclusions": "Because gravitational lens time delays are determined by the Hubble constant and the average surface density $\\kbar$ of the lens galaxy in the annulus between the images (Kochanek~\\cite{Kochanek02b}), we can make unambiguous estimates for the behavior of time delays in standard CDM halo models.  In these models, the expected delay is controlled by the mass fraction, $f_b=M_H/M_{vir}=\\Omega_{b,cold}/\\Omega_0$, in cold baryons making up the observed lens galaxy relative to the overall halo.  As the cold baryon fraction rises, so does the Hubble constant.  When the cold baryon fraction is comparable to the local baryonic content of galaxies ($f_b \\simeq 0.02$, Fukugita et al.~\\cite{Fukugita98}), the model parameters closely match those for isothermal (flat rotation curve) dark matter dominated lens models and the halos produce weak lensing signals compatible with weak lensing measurements in the SDSS (McKay et al.~\\cite{McKay01}). The mean surface density in the annulus is almost exactly $\\kbar=1/2$ and the local slope of the surface density is almost exactly $\\eta=2$ ($\\kappa \\propto R^{1-\\eta}$). Isothermal models are not only the best observational estimate for the lensing potential on these scales, they are also the model predicted by CDM assuming standard parameters and baryonic populations. For baryon fractions with a lower limit set by the local inventory and the upper limit set by the weak lensing measurements, we find that $H_0=48\\pm5$~km/s~Mpc based on the time delays measured for PG1115+080, SBS1520+530, B1600+434 and HE2149--2745. If all baryons were to cool and $f_b \\simeq 0.15\\pm0.05$, based on constraints from either the CMB (e.g. Netterfield et al.~\\cite{Netterfield02}, Wang et al.~\\cite{Wang02}) or cluster baryon fractions (e.g. White et al.~\\cite{White93}, Allen et al.~\\cite{Allen02}), then the Hubble constant could be as high as $H_0=65\\pm 6$~km/s~Mpc. Such models require most of the cold baryons in the lens galaxies to be in a locally invisible population and correspond to mass distributions less consistent with direct estimates. Both of these possibilities are lower than the local estimates of $H_0=72\\pm8$~km/s~Mpc from the HST Key Project (Freedman et al.~\\cite{Freedman01}), which agrees with the time delays of these four lenses only for mass distributions with constant $M/L$ ratios.  Thus, our detailed models for the expected properties of time delays in standard CDM halos agree with our simple models in Kochanek~(\\cite{Kochanek02a}), and we are faced with a conflict between CDM halo models, gravitational lens time delays and the local distance scale.  While there is some room for error in the lens results, the mutual agreement of the four simple, well-characterized time delay lenses and the simple relation between time delays, surface densities and the Hubble constant makes it difficult to point to a weakness (Kochanek~\\cite{Kochanek02b}). The most important observational steps are to improve the accuracies of the existing delay measurements and to expand the number of systems with delay measurements.  If the homogeneity of the results for simple lens systems, as compared to more complicated systems in clusters or with interacting galaxies, continues, the case for the existence of a conflict will become overwhelming. Improved characterizations of the lenses, either to constrain the mass distribution in the time delay lenses directly or to allow us to include the five other time delay lenses, are also important, but depend on obtaining deeper HST imaging of the systems.  Other constraints on the mass distributions such as weak lensing or the stellar dynamical measurements of the lens galaxies can also help to break any degeneracies.  In particular, estimates of the weak lensing signal as a function of the stellar velocity dispersion rather than luminosity would be excellent constraints on the halo extent in time delay lenses.  The systematic uncertainties in the mass distribution can be minimized by measuring time delays in lenses where the baryons dominate the mass and there there is little difference between a constant $M/L$ model and a model with dark matter.  This means measuring the time delay in very low redshift lens galaxies where the ratio of the critical radius to the effective radius, $R_c/R_e \\propto D_{OL}$ is small and the mass near the Einstein ring is increasingly dominated by the baryons.  For example, models of Q2237+0305 at $z_l=0.04$ suggest that less than 10\\% of the mass inside the Einstein ring of the lens can be dark, instead of the roughly 50\\% for typical models of higher redshift lenses (Trott \\& Webster~\\cite{Trott02}).  Unfortunately, Q2337+0305 has shown no variability on the very short time scale of its expected delay, making it a poor candidate for measuring time delays. There is a certain irony to proposing that local galaxies, which might be incorporated in local distance scale studies, are the ideal time delay lenses, but it may also lead to a system where the local and the ``cosmological'' distance scales can be compared directly.   \\noindent Acknowledgments. CSK thanks D. Rusin, P. Schechter, U. Seljak, J. Winn and S. Wyithe for discussions and comments.  CSK is supported by the Smithsonian Institution and NASA ATP grant NAG5-9265.  \\appendix"
    },
    "0206/astro-ph0206230_arXiv.txt": {
        "abstract": "We have proposed a viscoelastic model of the Maxwell stresses due to the disorganized magnetic field in MRI-driven MHD turbulence. Viscoelastic fluids in the laboratory are known to produce jet-like structures under the action of a rotating sphere. Here we argue that a similar mechanism may help explain jets in protostellar systems. Such jets would be driven not by large-scale organized magnetic fields, but by the mean-field stresses of small-scale tangled magnetic fields. ",
        "introduction": "Broadly speaking, the theory of jet acceleration and collimation has, over the past decades, progressed from hydrodynamical models to models incorporating magnetic fields. These models typically involve acceleration by what may be called ordered, large-scale fields. The long-range effects of ordered fields enable, for example, the coupling of a disk wind or jet to an accretion disk, which is a convenient source of power.  In contrast, the predominant role of magnetic fields in accretion disks is presumed to be the creation of a large effective turbulent viscosity, through the action of the Balbus-Hawley magnetorotational instability (MRI). The field produced by the MRI, as shown in simulations, has a significant contribution due to what might be characterized as a disorganized, or tangled, field. That a tangled field can have such important dynamical consequences as angular momentum transport in accretion disks leads one to ask what other consequences such a field might have. We propose that one consequence of this tangled field may be the driving of an axial outflow, {\\em i.e.} a jet. ",
        "conclusions": ""
    },
    "0206/astro-ph0206410_arXiv.txt": {
        "abstract": "{ We present two-dimensional (2D) radiative transfer modeling of {\\irc} at selected moments of its evolution in 1995--2001, which correspond to three epochs of our series of 8 near-infrared speckle images \\citep{Osterbart_etal2000, Weigelt_etal2002}. The high-resolution images obtained over the last 5.4 years revealed the dynamic evolution of the subarcsecond dusty environment of {\\irc} and our recent time-independent 2D radiative transfer modeling reconstructed its physical properties at the single epoch of January 1997 \\citep{Men'shchikov_etal2001}. Having documented the complex changes in the innermost bipolar shell of the carbon star, we incorporate the evolutionary constraints into our new modeling to understand the physical reasons for the observed changes. The new calculations show that our previous static model is consistent with the brightness variations seen in the near-infrared images, implying that during the last 50 years, we have been witnessing an episode of a steadily increasing mass loss from the central star, from $\\dot{M} \\approx 10^{-5}$ {\\Msun}\\,yr$^{-1}$ to the rate of $\\dot{M} \\approx 3 \\times 10^{-4}$ {\\Msun}\\,yr$^{-1}$ in 2001. The rapid increase of the mass loss of {\\irc} and continuing time-dependent dust formation and destruction caused the observed displacement of the initially faint components C and D and of the bright cavity A from the star which has almost disappeared in our images in 2001. Increasing dust optical depths are causing strong backwarming that leads to higher temperatures in the dust formation zone, displacing the latter outward with a velocity $v_T \\approx 27$ km\\,s$^{-1}$ due to the evaporation of the recently formed dust grains. This self-regulating shift of the dust density peak in the bipolar shell mimics a rapid radial expansion, whereas the actual outflow has probably a lower speed $v < v_\\infty \\approx 15$ km\\,s$^{-1}$. The model predicts that the star will remain obscured until $\\dot{M}$ starts to drop back to lower values in the dust formation zone; in a few years from that moment, we could be witnessing the star reappearing. \\keywords { radiative transfer -- circumstellar matter -- stars: individual: {\\irc} -- stars: mass-loss -- stars: AGB and post-AGB -- infrared: stars } } ",
        "introduction": "\\label{Introduction}  The pulsating carbon star {\\irc} (also known as \\object{CW Leo}, \\object{AFGL 1381}), together with its huge circumstellar envelope lost during its long evolution on the asymptotic giant branch (AGB), is the best studied object of its class. Being in a very advanced phase of its life, probably in transition to protoplanetary nebulae \\citep[][ hereafter Paper~I]{Osterbart_etal2000}, {\\irc} presently exhibits a very high mass-loss rate $\\dot{M} \\sim 10^{-4}$ {\\Msun}\\,yr$^{-1}$ \\citep[][ hereafter Paper~II]{Men'shchikov_etal2001}. After three decades of intensive observational and theoretical work \\citepalias[see, e.g., references in][]{Men'shchikov_etal2001}, recent near-infrared speckle imaging has revealed an extremely complex evolution of its circumstellar material in the vicinity of the dust condensation zone on a time scale of one year.  Near-infrared images with resolutions better than 100 mas presented by \\citep{Weigelt_etal1997,Osterbart_etal1997,Weigelt_etal1998a,Weigelt_etal1998b, HaniffBuscher1998, Osterbart_etal2000,Tuthill_etal2000,Weigelt_etal2002} have demonstrated that the inner, subarcsecond dust shell of {\\irc} is non-spherical and clumpy, with four components A, B, C, and D clearly visible. The detailed two-dimensional radiative transfer modeling presented in \\citetalias{Men'shchikov_etal2001} has shown that the star is actually located at the position of the second brightest component B. The brightest southern peak A was identified with the radiation emitted and scattered in the optically thinner cavity of the dense circumstellar shell. The model reconstructed physical properties of the star and dusty envelope of {\\irc} at a single moment corresponding to the epoch of our high-resolution $H$- and $K$-band images on January 23, 1997 \\citepalias{Osterbart_etal2000}. As much as it was possible, the model took into account most other observations of dust radiation, although stellar pulsations and non-periodic changes of the shell made many measurements from various epochs fundamentally incomparable in the frame of the static model.  In the present study, we attempt to attack the problem using a simplified approach based on the self-consistent model for a single epoch that we have constructed in \\citetalias{Men'shchikov_etal2001}. The idea is to extend this modeling to the first and to the last epochs of the 6-year sequence of high-resolution $K$-band speckle images we have obtained since October 1995 \\citep[][ hereafter Paper~III]{Weigelt_etal2002}. In the beginning of our monitoring of {\\irc}, component B (the star) was relatively bright. The direct light from the star has been gradually fading since then, whereas the bright lobe A (the southern cavity) has become dominant. The angular distance between components A and B increased from $\\sim$ 190 to $\\sim$ 350 mas between 1995 and 2001, implying a linear speed $v_{\\rm A} \\approx 18$\\,km\\,s$^{-1}$ in the plane of sky and a deprojected radial velocity $v_{r,{\\rm A}} \\approx 19$\\,km\\,s$^{-1}$ \\citepalias[Appendix A of][]{Men'shchikov_etal2001}, for the assumed distance of $D = 130$ pc and the viewing angle of $\\theta_{\\rm v} = 40${\\degr} \\citepalias{Men'shchikov_etal2001}. On the basis of our time-independent model, this well-documented rapid evolution has been qualitatively interpreted in terms of dust formation in a bipolar stellar wind, the increasing mass-loss rate, and the sublimation of the recently formed grains in a progressively hotter region just outside the dust formation radius \\citepalias{Weigelt_etal2002}.  The goal of the present study is to determine, whether the observed changes in the high-resolution speckle images are consistent with our previous model of {\\irc}, and to derive more accurate physical parameters of the wind at several moments in time. In Sect.~\\ref{RadTraModel}, we describe our approach and model parameters. In Sect.~\\ref{Results}, we discuss the results of our new modeling, comparing them with the high-resolution images, intensity profiles, and spectral energy distribution (SED) of {\\irc}. In Sect.~\\ref{Conclusions}, we summarize the model parameters and our conclusions. ",
        "conclusions": "\\label{Conclusions}  We presented results of two-dimensional radiative transfer modeling of {\\irc} at selected moments of its recent evolution, which correspond to three epochs of our series of high-resolution speckle images of the object recorded over the last 5.4 years \\citepalias{Osterbart_etal2000,Weigelt_etal2002}. The imaging revealed dynamic evolution of the subarcsecond dusty environment of {\\irc} on a very short time scale, of the order of one year. The goal of the present study was to test predictions of our recent time-independent modeling \\citepalias{Men'shchikov_etal2001} which derived the structure and physical properties of {\\irc} at a single epoch (January 1997). Now that the entire sequence of 8 high-resolution near-IR images has documented complex changes in the inner bipolar dusty shell of the carbon star, we took these temporal constraints into account to see, whether the static model is consistent with the fast {\\em evolution} seen in our images. The new modeling allowed us to make a quantitative physical interpretation of what goes on in {\\irc}.  Our previous static model of {\\irc} \\citepalias{Men'shchikov_etal2001} is consistent with the new constraints, provided that small modifications are made to its density profile in the dust formation zone and to the opening angle of the bipolar cavity. The new modeling has shown that the cavity has been shrinking from $\\omega_1 \\approx 36${\\degr} to $\\omega_8 \\approx 26${\\degr} and the density distribution across it has been changing during the 5.4 years of our imaging. We are witnessing a dynamic episode of rising mass loss from the central star, from $\\dot{M} \\approx 10^{-5}$ {\\Msun}\\,yr$^{-1}$ to the present value of $\\dot{M} \\approx 3 \\times 10^{-4}$ {\\Msun}\\,yr$^{-1}$, which probably started $\\sim$ 50 years ago. If the current rate of the increase of $\\dot{M}$ is constant, the mass-loss rate from the stellar surface may be now as high as $2.6 \\times 10^{-3}$ {\\Msun}\\,yr$^{-1}$.  A compact dense shell with bipolar cavities has formed around the star as a result of the rapid increase of the mass loss by {\\irc}, causing the observed rapid evolution documented in our images. The higher mass loss produces favorable conditions for dust formation in the increasingly dense inner envelope expanding outward with the outflow velocity $v \\approx 15$ km\\,s$^{-1}$. Larger amounts of dust increased the optical depths, obscuring the central star, whereas the optically thinner cavity remained relatively unaffected. Due to backwarming, temperatures in the dust formation zone became higher, shifting the latter to larger radii at the velocity $v_T \\approx 27$ km\\,s$^{-1}$ and mimicking the real outflow motion of the circumstellar shell material with that speed. One can predict that the star will remain obscured until $\\dot{M}$ starts to drop back to lower values in the dust formation zone. Within a few years from that moment, we could be witnessing the star reappearing, whereas the cavities becoming relatively fainter."
    },
    "0206/astro-ph0206357_arXiv.txt": {
        "abstract": "Two approximations, namely the \\SA\\ and the \\AAP, are presently used to filter out the acoustic modes when computing low frequency modes of a star (gravity modes or inertial modes). In a precedent paper (Dintrans \\& Rieutord 2001), we observed that the \\AAP\\ gave eigenfrequencies much closer to the exact ones than the \\SA. Here, we try to clarify this behaviour and show that it is due to the different physical approach taken by each approximation: On the one hand, the \\SA\\ considers the low frequency part of the spectrum of (say) gravity modes and turns out to be valid only in the central region of a star; on the other hand, the \\AAP\\ considers the \\BV frequency as asymptotically small and makes no assumption on the order of the modes. Both approximations fail to describe the modes in the surface layers but eigenmodes issued from the \\AAP\\ are closer to those including acoustic effects than their subseismic equivalent.  We conclude that, as far as stellar eigenvalue problems are concerned, the \\AAP\\ is better suited for simplifying the eigenvalue problem when low-frequency modes of a star are considered, while the \\SA\\ is a useful concept when analytic solutions of high order low-frequency modes are needed in the central region of a star. ",
        "introduction": "When considering the low-frequency modes of a star, namely gravity modes or inertial modes, the compressibility of the fluid is often a side effect in the determination of eigenfrequencies and eigenmodes; in other words, the dynamics of these modes may be simplified by neglecting the elasticity of the fluid or, equivalently, by filtering out acoustic modes. This is the aim of the subseismic and anelastic approximations; the resulting equations for eigenmodes are much simpler than the original ones and very useful when dealing with the low frequency oscillations of rotating stars \\cite[e.g.][]{DR00}.  Recently, we compared these two approximations (\\cite{DR01} referred to as paper I hereafter). We found that in the two cases which we analysed, namely two polytropes, the \\AAP\\ performed much better than the \\SA. We attributed this behaviour to an inconsistency of the \\SA\\ but our argument turns out to be not general and \\cite{Smeyers01} showed that, for low-frequency high order modes, the \\SA\\ gives the first order equations in regions not close to the surface of the star. These results prompted us to re-examine this question in order to clarify the origin of the different behaviour of the two approximations. For this purpose we will focus, in section 2, on two asymptotic developments: a first one where we use, as \\cite{Smeyers01}, the frequency as a small parameter and a second one where we use the \\BV frequency as the small parameter. These asymptotic developments will prove to be at the origin of each of these approximations and will allow us to clarify the physics attached to each of them. In section 3, using the same examples as in paper I, we will compare the approximate eigenfunctions to their exact counterparts and show the better behaviour of the \\AAP. Our conclusions are drawn in section 4. ",
        "conclusions": "\\psfrag{0}{0}\\psfrag{omegan}{$\\omega_N$}\\psfrag{omegac}{$\\omega_c$} \\psfrag{omega}{$\\omega$} \\psfrag{eps}{$\\eps\\omega$} \\psfrag{gravity}{\\hspace*{-7.5mm}$\\overbrace{\\hspace*{2.5cm}}^{\\rm \\displaystyle Gravity \\; modes}$} \\psfrag{acoustic}{\\hspace*{-9.5mm}$\\overbrace{\\hspace*{2cm}}^{ \\parbox{2.5cm}{\\rm Acoustic modes\\\\ rejected to $\\infty$}}$}  \\begin{figure} \\vspace*{1cm} \\centerline{\\includegraphics[width=8cm]{an_spectrum.eps}} \\vspace*{1cm} \\psfrag{gravity}{\\hspace*{-1.0cm}$\\overbrace{\\hspace*{1.2cm}}^{ \\parbox{2.5cm}{Asymptotic gravity\\\\ \\centerline{modes}}}$} \\psfrag{acoustic}{\\hspace*{-1.2cm}$\\overbrace{\\hspace*{2.5cm}}^{ \\parbox{2.5cm}{\\rm Acoustic modes}}$} \\centerline{\\includegraphics[width=8cm]{sa_spectrum.eps}} \\caption[]{A schematic picture of the modes of a star viewed from the anelastic viewpoint (above) and subseismic viewpoint (below). $\\omega_N$ and $\\omega_c$ are respectively the frequency of the lowest order gravity and acoustic modes.} \\label{schema} \\end{figure}  In this paper we tried to clarify the differences between the subseismic and anelastic approximations which both aim at describing the low frequency spectrum. The \\SA\\ appears when one concentrates on the low frequency high radial order modes in the central region of a star; no constraint is imposed to the \\BV frequency.  On the other hand the \\AAP\\ assumes a weak stratification but imposes no constraint on the degree of the mode.  Hence, while the \\AAP\\ makes the \\BV frequency, and thus the frequency of all gravity modes, vanishingly small compared to acoustic frequencies, the \\SA\\ focuses on gravity modes whose radial order is very large and hence have small frequencies compared to acoustic ones.  In other words, the \\AAP\\ removes the elasticity of the fluid by rejecting acoustic frequencies to infinity and therefore allows for a description of the full spectrum of gravity modes while the \\SA, keeping $\\omega_c$ and $\\omega_N$ in a finite ratio, concentrates on one part of the spectrum, namely that containing high radial order modes which are the least sensitive to the elasticity of the fluid.  This situation is summarized in figure~\\ref{schema}.  Since in stars the situation is often that $\\omega_{N}\\ll \\omega_{c}$, the use of the \\AAP\\ is recommended as it is likely closer to the solutions of the complete equations; on the other hand, the \\SA\\ may be useful when one needs an analytic expression of gravity modes in the central regions of a star.  Finally, it is worth mentioning the work of \\cite{Durr89} who discussed these two approximations in the context of atmospheric sciences. In this field, where the \\SA\\ is called the ``pseudo-incompressible appoximation\" and the anelastic approximation the ``modified anelastic approximation\", the \\SA\\ appears to be superior to the \\AAP\\ as it conserves the energy, a property which is important for nonlinear problems.  This result shows that the best choice for filtering out acoustic modes is dependent on the problem at hands.  Therefore, our results which favour the \\AAP\\ when searching for low-frequency modes of stars, may be specific to eigenvalue problems."
    },
    "0206/astro-ph0206161_arXiv.txt": {
        "abstract": "In the standard Cold Dark Matter model of structure formation, massive clusters form via the merger of smaller clusters. N-body/hydrodynamical simulations of merging galaxy clusters have shown that mergers can temporarily boost the X-ray luminosity and temperature of the merged cluster above the equilibrium values for the merged system. The cumulative effect of these ``merger boosts'' will affect the observed X-ray luminosity functions (XLFs) and temperature functions (TFs) of clusters. One expects this effect to be most important for the most luminous and hottest clusters. XLFs and TFs of clusters provide some of the strongest constraints on cosmological and large-scale-structure parameters, such as the mean fluctuation parameter, $\\sigma_8$, and the matter density divided by the critical density, $\\Omega_0$. Merger boosts may bias the values of $\\sigma_8$ and $\\Omega_0$ inferred from cluster XLFs and TFs if virial equilibrium is assumed. We use a semi-analytic technique to estimate the effect of merger boosts on the X-ray luminosity and temperature functions. The boosts from individual mergers are derived from N-body/hydrodynamical simulations of mergers. The statistics of the merger histories of clusters are determined from extended Press-Schechter (PS) merger trees. We find that merger boosts can increase the apparent number of hot, luminous clusters. For example, in a Universe with $\\Omega_0 = 0.3$ and $\\Omega_\\Lambda = 0.7$ at a redshift of $z=1$, the number of clusters with temperatures $T > 10$ keV is increased by a factor of 9.5, and the number of clusters with luminosities $L_X > 5 \\times 10^{44} \\, h^{-2}$ erg s$^{-1}$ is increased by a factor of 8.9. We have used our merger-boosted TFs and XLFs to derive the cosmological structure parameters $\\sigma_8$ and $\\Omega_0$ by fitting Press-Schechter equilibrium relations to local ($z=0$) and distant (either $z=0.5$ or $z=1$) cluster samples. Merger boosts cause $\\sigma_8$ to be overestimated by about 20\\%. The matter density parameter $\\Omega_0$ may be underestimated by about 20\\%, although this result is less clear. If the parameters of the fluctuation spectrum are derived from the observed TF or XLF (e.g., from a low redshift sample), then this removes most of the boost effect on $\\Omega_0$. However, larger merger boost effects may appear when cluster XLFs and TFs are compared to cosmological structure parameters derived by other techniques (e.g., cosmic microwave background fluctuations or the brightness of distant supernovae). ",
        "introduction": "\\label{sec:intro}  Clusters of galaxies have been widely used to provide useful constraints on cosmological parameters (e.g., Henry \\& Arnaud 1991; Bahcall \\& Fan 1998; Eke et al.\\ 1998; Borgani et al.\\ 2001). This is in part possible because there is a well-developed theoretical framework which allows one to predict the mass function (MF) of clusters of galaxies and its evolution as a function of the cosmology. Here, the MF is defined as the number density of clusters as a function of their mass. One standard method for predicting the MF is to use Press-Schechter formalism, originally developed by Press \\& Schechter (1974, hereafter PS), and developed in more detail by Bond et al.\\ (1991) and Lacey \\& Cole (1993), among others, in combination with the Cold Dark Matter (CDM) model. In hierarchical structure formation models like CDM, more massive halos form from the merger of smaller halos. Values of cosmological parameters can be estimated by comparing theoretical models for the MF of clusters with observations.  This technique places the strongest constraints on $\\sigma_8$, the RMS mass fluctuations on a scale of $8 \\, h^{-1}$ Mpc where $h$ is the Hubble constant in units of $100$ km/sec/Mpc, and on $\\Omega_{0} \\equiv \\bar{\\rho} / {\\rho_c}$, the ratio of the current mean matter density to the critical density $\\rho_c = 3 H_0^2 / ( 8 \\pi G )$ at the current epoch (e.g., Henry \\& Arnaud 1991; Kitayama \\& Suto 1996; Eke et al.\\ 1998; Borgani et al.\\ 2001; Ikebe et al.\\ 2002). Roughly speaking, the present-day abundance of clusters determines the relationship between $\\sigma_8$ and $\\Omega_{0}$, while the evolution of clustering with redshift breaks this degeneracy, although Reiprich \\& B\\\"ohringer (2001) have been able to constrain $\\Omega_{0}$ using only a local cluster sample by assuming a CDM spectrum of perturbations whose shape parameter depends on $\\Omega_{0}$. In these determinations, the most massive clusters have the greatest leverage. Massive clusters are rare objects, and the abundance of the most massive clusters is very sensitive to the cosmological parameters.  Unfortunately, the MF of clusters of galaxies cannot be directly observed (save by gravitational lensing in a relatively small number of cases). Generally, it is inferred from the X-ray luminosity function (XLF) or temperature function (TF), using empirical or semi-empirical relations between the masses of clusters and their temperatures ($T$) or X-ray luminosities ($L_X$). These scaling relations are usually applied under the assumption that the clusters are dynamically relaxed, although in reality this cannot always be true since in the CDM model larger clusters are continually forming via mergers of smaller clusters.  Simulations have shown that if two clusters of comparable mass merge to form a larger cluster, there is a temporary increase in the cluster's X-ray luminosity and temperature (Ricker \\& Sarazin 2001; Ritchie \\& Thomas 2002). If the cluster is observed during this period of boosted luminosity and temperature, the inferred mass will be larger than the actual mass of the cluster. In fact, such objects will be preferentially detected in X-ray flux-limited samples because they are intrinsically more luminous than equal mass clusters which are not currently experiencing a merger. This will be particularly true for X-ray selected, high redshift clusters. As a possible examples of this, one of the most distant X-ray selected clusters observed to date, RXJ1053.7+5735, has a morphology which may indicate an ongoing merger (Hashimoto et al.\\ 2002), as does the most luminous X-ray selected cluster found to date, RXJ1347.5-1145 (Allen, Schmidt, \\& Fabian 2002).  If this X-ray luminosity-temperature (hereafter L-T) boost effect is sufficiently large, and if mergers occur frequently enough, then the observed XLF and TF, and hence the inferred MF, will be different from what they would be if all clusters were dynamically relaxed. One would expect that merger boosts would affect most strongly the high luminosity and temperature ends of the XLF and TF. Since the massive clusters which have very high values of $L_X$ and $T$ are rare, a small contribution of lower mass clusters with L-T boosts could strongly affect the statistics. As we noted above, the abundance of the hottest and  most luminous clusters (which are usually assumed to be the most massive clusters) has a strong influence on the inferred values of the cosmological parameters. In particular, if merger boosts artificially increase the abundance of hot, luminous clusters at high redshift, the real value of $\\Omega_{0}$ will be underestimated. Since observations of the numbers of high temperature and luminosity clusters at moderate and high redshifts have been used to infer that $\\Omega_{0}$ is low (e.g., Bahcall, Fan, \\& Cen 1997; Borgani et al.\\ 2001), this effect could be important.  In this paper we attempt to quantify the effect of mergers on the observed XLF and TF, and to determine how the X-ray L-T boost associated with mergers alters the values of $\\sigma_8$ and $\\Omega_{0}$ inferred from observations. We consider three possible cosmologies: an ``open'' model ($\\Omega_{0} = 0.3$, $\\Omega_\\Lambda = 0$), a ``flat'' model ($\\Omega_{0} = 0.3, \\Omega_{\\Lambda} = 0.7$), and an Einstein-de Sitter (EdS) model ($\\Omega_{0} = 1$, $\\Omega_\\Lambda = 0$). Here, $\\Omega_\\Lambda = \\Lambda/3 H_0^2$ and $\\Lambda$ is the cosmological constant. We use Extended Press-Schechter (EPS) theory (\\S~\\ref{sec:basic_ps}) and a Monte Carlo technique to build a collection of merger histories of clusters (``merger trees,'' \\S~\\ref{sec:trees}) for a variety of cosmologies. We then build two sets of XLFs and TFs from the MF given by the merger trees for each cosmology by averaging together many merger trees (\\S~\\ref{sec:XLF+TF}). One set assumes that clusters are always relaxed when transforming from mass to temperature or luminosity (\\S\\S~\\ref{sec:M-T+M-L}, \\ref{sec:no_boosts}), while the other set includes the L-T boost effect (\\S~\\ref{sec:with_boosts}). We ignore such non-gravitational effects as preheating since such effects are mainly important for lower mass clusters and groups (Ponman et al.\\ 1996), whereas we choose to focus on massive clusters where the effects of merger boosts are most important. We also ignore cooling flows at the centers of clusters, which may be disrupted by mergers. The disruption of cooling flows may increase the effect of mergers on X-ray temperatures, but decrease their effect on X-ray luminosities. The strength and duration of the L-T boost is estimated for arbitrary merging masses and impact parameters by interpolating and extrapolating from the results of a series of N-body/hydrodynamical simulations of binary cluster mergers (\\S~\\ref{sec:hydro} and Appendix~B). The angular momenta or impact parameters for the mergers are drawn from a distribution based on linear theory (\\S~\\ref{sec:impact_param}). The boosted XLFs and TFs are presented in \\S~\\ref{sec:XLF+TF} at a variety of redshifts, and are compared to unboosted results. The boosted and unboosted XLFs and TFs at several redshifts are fitted using a PS mass function and equilibrium relations between the mass and temperature or X-ray luminosity (\\S~\\ref{sec:tree_fit}). The best-fit parameters, specifically $\\sigma_8$ and $\\Omega_{0}$, are compared to the ``actual'' parameters used to build the trees. Our conclusions are summarized in \\S~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  Hydrodynamical simulations of binary cluster mergers have shown that mergers can temporarily boost the X-ray luminosity and temperature of the merged cluster beyond their virial equilibrium values. These merger boosts can alter the observed TFs and XLFs of clusters, particularly at the high ends. Even a few ``extra'' clusters with high luminosities and temperatures can have a significant impact due to the relative rarity of massive clusters. The mass function of clusters is commonly inferred from the observed TF or XLF, assuming virial equilibrium. The inferred mass function is often used to constrain cosmological parameters. Thus, merger boosts can affect the inferred cosmological parameters.  We tested the amplitude of the merger boost effect by using EPS theory and a Monte Carlo technique to numerically reconstruct the merger histories (trees) of a population of clusters. At each merger, we quantified the strength of the merger boost by extrapolating or interpolating from a set of hydrodynamical simulations of cluster mergers. We then built TFs and XLFs from the merger trees, averaging together our sets of merger histories so that the sample cluster population was representative of the observed present day population.  Our results show that merger boosts do indeed affect the apparent number of high mass clusters. For example, in the flat model at $z=1$ merger boosts cause the number of clusters with temperatures $T > 10$ keV to increase by a factor of 9.5, and the number of clusters with luminosities $L_X > 5 \\times 10^{44} \\, h^{-2}$ erg s$^{-1}$ is increased by a factor of 8.9.  The effect is strongest for an EdS model, since clusters evolve more rapidly in this model than in the open and flat models.  In the EdS model at a redshift of $z=0.5$, the number of clusters with temperatures $T > 10$ keV is increased by a factor of 38, and the number with luminosities $L_X > 5 \\times 10^{44} \\, h^{-2}$ erg s$^{-1}$ by a factor of 24. At first this might appear contradictory with the fact that the X-ray luminosity receives a larger boost from mergers than the X-ray temperature (see Figure~\\ref{fig:LtTt}).  However, the range of X-ray luminosities is larger than the range of X-ray temperatures; another way of saying this is that the X-ray luminosity-temperature relationship is much steeper than linear. At a result, the overall effect of luminosity boosts on the XLF is not as strong as the effect of temperature boosts on the TF.  Comparing the boosted and unboosted differential TF for the EdS and flat models at $z=1$ shows that merger boosts cause $T \\approx 10$ keV clusters to be almost as common in the EdS model as in the flat model.  This means that if all cosmological parameters other than $\\Omega_0$ were known from other measurements, and the abundance of very hot or very X-ray luminous clusters at high redshift were used to determine $\\Omega_0$, the value would be substantially underestimated.  We fit PS distributions to our simulated TFs and XLFs. We did this by considering two samples: a low redshift, local sample ($z = 0$), and a moderate or high redshift sample ($z = 0.5$ or $z = 1$). The two samples are fit simultaneously to determine the best-fit values of $\\sigma_8$ and $\\Omega_0$, and these values were compared to the actual values used to construct the merger trees. Merger boosts can cause $\\sigma_8$ to be overestimated by about 20\\%. Merger boosts may also cause $\\Omega_0$ to be underestimated by about 20\\%, although the results from the XLF fits do not show a clear trend. The effect of merger boosts on the inferred value of $\\Omega_0$ is much smaller in these fits than might have been expected from the overall increase in the number of hot clusters at high redshifts produced by merger boosts. Much of the effect of merger boost is ``renormalized'' away by the joint fit of low and high redshift sample. One way to think of this is that merger boosts increase the numbers of hot clusters, both at low redshift and high redshift. Determining $\\sigma_8$ from the low redshift sample or from a joint fit removes much of the boost effect, and the change in $\\Omega_0$ is smaller than one might have expected.   Studies of the abundance of hot or luminous clusters at high redshift have been used to argue that we live in a low density Universe (e.g., Bahcall \\& Fan 1998; Borgani et al.\\ 2001). It would appear that merger boosts do not invalidate this conclusion, although the error bars should be increased significantly to include the systematic uncertainties associated with merger boosts.  On the other hand, merger boosts do affect the value of $\\sigma_8$. As noted before, the number of the hottest and most luminous clusters are affected quite strongly. Thus, any comparison between cluster TFs and XLFs and cosmological parameters derived from other objects (from cosmic microwave background radiation fluctuations, or from the brightness of distant supernovae) is likely to be inconsistent and may lead to errors in the deduced cosmological parameters. At the least, the systematic uncertainties are likely to be much larger than might be inferred from the statistics of clusters alone.  Obviously, the effect of merger boosts could be avoided entirely if the mass function could be determined directly from gravitational lensing observations. Radio detections of the Sunyaev-Zel'dovich (SZ) effect can also provide a measure of the number of luminous and hot clusters, particularly at high redshift (e.g., Holder et al.\\ 2000). Although mergers should also boost the microwave decrement from clusters, we expect that this effect would be smaller than the effect on the X-ray emission-weighted temperature or the X-ray luminosity, because the SZ effect depends on density rather than the square of the density."
    },
    "0206/astro-ph0206482_arXiv.txt": {
        "abstract": "% We summarize the results of an observational program performed with the satellite \\sax with the aim to find and study more extreme BL Lac objects. We  discuss the SEDs of the observed objects and their impact on the ``blazar sequence\" scenario, and consider their relevance as possible TeV emitting sources. ",
        "introduction": "One of the main differences among  blazars is constituted by the  position of the peak of the synchrotron component in their spectral energy distribution (SED), namely at low (mm--IR) or high (UV and soft X--ray) frequencies (e.g. Padovani \\& Urry 2001). The results of the \\sax and ASCA observations  have shown that there is a rather smooth sequence with respect to the peak frequencies, and in particular that this sequence is well extended also at high energies ($>1$ keV), in a range of physical conditions not previously considered. With this respect, the \\sax observations of Mkn 501 and 1ES 2344+514 have been fundamental, revealing for the first time objects with synchrotron peak frequencies around or above 100 keV. Such sources are of great interest also because some of them have been detected at TeV energies by Cherenkov telescopes (Catanese \\& Weekes 1999). In these objects, X-ray and TeV 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 strong correlation between TeV and X-ray emissions, clearly evident in the 1997 flare of Mkn 501 (Pian et al. 1998, Aharonian et al. 1997) and in Mkn 421 (Maraschi et al. 1999, Krawczynski et al. 2001), together with the very rapid variability displayed (Mkn 421 doubled its TeV flux in less than 20 min, Gaidos et al. 1996), provides very strong constrains for any emission model and a powerful tool for diagnostics. TeV BL Lacs are also interesting because, being the only known extragalactic sources at these energies, they allow an independent estimate of the extragalactic IR background (IRB), due to the absorption of high energy photons through $\\gamma-\\gamma$ collision and pair production (Stecker et al. 1992).     With the aim to find and study more ``extreme\"\\footnote{we will call ``extreme\" a source with $\\nu_{\\rm peak}\\gsim1$ keV} objects, in order to sample more accurately the high energy branch of the peak sequence, we have performed an observational campaign with \\sax, taking advantage of its unique  wide energy band (0.1-100 keV). Here we summarize the  results  on 9 sources, observed  between 1998 and 2001, in the context of their SED. We then consider the  impact of these objects in the determination of the intrinsic physical parameters governing the peak sequence, and present a selection criterium according to which they can be considered good targets for Cherenkov telescopes. ",
        "conclusions": ""
    },
    "0206/astro-ph0206211_arXiv.txt": {
        "abstract": "We compare the diffraction-limited field of view (FOV) provided by four types of off-axis Gregorian telescopes: the classical Gregorian, the aplanatic Gregorian, and designs that cancel astigmatism and both astigmatism and coma. The analysis is carried out using telescope parameters that are appropriate for satellite and balloon-borne millimeter and sub-millimeter wave astrophysics.  We find that the design that cancels both coma and astigmatism provides the largest flat FOV, about 21 square degrees.  We also find that the FOV can be increased by about 15$\\%$ by optimizing the shape and location of the focal surface. ",
        "introduction": "\\label{Intro}  Future advances in mm and sub-mm wave astronomy critically depend on the design of optical systems. Detector technology has matured to the point where detector sensitivity is limited by photon noise from the source or from unavoidable photons in the light-path, such as the atmosphere and the mirrors.  Only an increase in the number of photometers can significantly increase the overall sensitivity of an instrument.  Arrays of detectors with tens and hundreds of elements have recently come on line\\cite{scuba,bolocam} and it is widely expected that the construction of such instruments will accelerate as the fabrication of both semiconductor-based and superconducting-based bolometers becomes more uniform and more automated through the use of standard micro-lithography techniques \\cite{Agnese99,Benford98,Bock98,Kreysa98,Gildemeister00}.  These large focal plane arrays need to be coupled to telescopes that provide a correspondingly large, diffraction-limited field of view (FOV). The need for optical systems with large useable FOV was not acute in the past when typically only a few photometers were coupled to the telescope. Because of the long wavelength, it is relatively easy to design an optical system for use at the millimeter wave band that is diffraction limited near the center of the FOV. The current challenge is to provide for the {\\it largest} useable FOV in order to accommodate large arrays.  An increase in the available FOV of off-axis Gregorian telescopes is of particular interest because such telescopes have higher aperture efficiency and lower side-lobe response than Cassegrain or on-axis Gregorian reflecting telescopes \\cite{DragoneHogg}.  It is also interesting to analyze in detail Gregorian telescopes that have a low f-number and a small number of mirrors because such systems find wide use in satellites and balloon borne payloads which require compact and simple optical systems.  For example, Gregorian telescopes have been used extensively in recent years in ground based and balloon borne experiments to characterize the anisotropy of the cosmic microwave background radiation (CMB) \\cite{archeops,maxima,boomerang,viper}. Both NASA's Microwave Anisotropy Probe (MAP) satellite and the European Space Agency's Planck satellite, two missions designed to map CMB temperature fluctuations, employ off-axis Gregorian telescopes with low f-number \\cite{map,planck}.  The primary source of aberrations in Gregorian and Cassegrain telescopes are coma and astigmatism. Several designs have been proposed to improve on the classical Gregorian (CG) design, which has a parabolic primary and an elliptical secondary.  In an aplanatic Gregorian (AG) telescope coma is cancelled without creating spherical aberration. The primary mirror is slightly ellipsoidal and the secondary is a slightly more eccentric ellipsoid than in the similar CG. The conic constant of one mirror is chosen to eliminate coma and the conic constant of the other mirror is adjusted to compensate the spherical aberration introduced by the change in the first mirror. Aplanatic designs are fairly common, for example, aplanatic versions of Cassegrain telescopes were used in the Hubble Space Telescope \\cite{hubble} and the Keck 10-meter telescopes \\cite{keck}.  Dragone \\cite{Dragone82,Dragone83} has described designs for off-axis Gregorian systems which eliminate astigmatism, and both astigmatism and coma; hereafter we refer to these designs as D1 and D2, respectively. These designs also greatly reduce instrumental polarization for field points near the center of the FOV \\cite{Mizuguchi78}. The reduction in instrumental polarization is of benefit for antennas for communication systems, which use polarization as a method to increase band-width \\cite{Dragone78,Westcott79}, and for experiments designed for detecting polarized signals. For example, intense efforts are now being made by a number of experiments to discover the CMB polarization anisotropy \\cite{polatron,pique,maxipol,polar,compass}.  The AG, D1 and D2 designs present progressively improved image quality near the center of the field of view, however it is not clear which of the systems provides a larger useable FOV, which is the quantity of interest for millimeter wavelength focal plane arrays.  In this paper we quantitatively compare the size of the diffraction-limited FOV provided by these three optical designs.  Because we are interested in potential applications for millimeter-wave astrophysics, and CMB research in particular, we perform our analysis with telescopes that provide $\\sim 8$ arcminute full-width at half-maximum beam size at 150 GHz ($ \\lambda = 2$ mm). We comment on the applicability of our analysis to other wavelengths in Section 3.  We use the telescope of the Archeops balloon borne experiment\\cite{archeops} as a baseline for comparison. Archeops is designed to observe the CMB with an array of 24 bolometric photometers distributed in four frequency bands between 143 and 545 GHz with beam sizes between 8 and 5 arcminutes, respectively. The focal plane array is a prototype of the High Frequency Instrument, one of two focal plane instruments on board ESA's Planck satellite. The satellite is scheduled to be launched around 2007. Archeops has an off-axis tilted Gregorian telescope following the D1 design, and is similar in its physical parameters to the Planck telescope \\cite{planck}. ",
        "conclusions": "\\label{Results}  The results, summarized in Table~\\ref{Table-FOV}, show that the D2 design provides the largest available FOV and is a good choice as a telescope that needs to accommodate a large array of photometers. The AG design provides the smallest FOV, although it is still considerably better than the classical Gregorian telescope upon which all of these systems are based.  With our telescope parameters, the optimized D1 and D2 designs provide $\\sim 20\\%$ and $\\sim 50\\%$ larger DLFOV than the optimized AG design, respectively. As expected, a larger DLFOV is obtained in all systems by optimizing the parameters of the focal surface, but this improvement decreases from $\\sim 60\\%$ in the case of the AG design to only about 20\\% for the D2 design. Because the D2 system with a flat focal plane provides a useable FOV that is almost as large as the one with an optimized focal surface, it is very suitable for arrays of detectors that are fabricated on flat silicon wafers \\cite{Agnese99,Benford98,Bock98,Kreysa98,Gildemeister00}. For this system the physical lengths of the axes of the diffraction limited region of the focal plane are 17.7 and 16.2 cm in the elevation and cross-elevation directions, respectively.  The DLFOV that we found for each of the three telescope designs is the area in which the effects of aberrations are small compared to diffraction for a frequency of 150 GHz and an aperture that gave a single mode beam size of 8 arcminutes. It is straight-forward to show that for a fixed beam size and in the single mode optics limit the DLFOV will be larger at higher frequencies.  For single mode optics the aperture area $A$, beam solid angle $\\Omega$, and frequency $\\nu$ are related through $A\\Omega = C/\\nu^{2}$, so for a fixed beam size $A \\propto \\nu^{-2}$; at higher frequencies the aperture area is smaller. Since the diffraction spot (or Airy disk) size scales as $1/(A \\nu^{2})$, it is a constant as a function of frequency under these assumptions. Ray aberrations, however, decrease as $A$ decreases so the ratio of aberration size to diffraction spot size also decreases with increasing $\\nu$.  Because a single mode system of constant beamsize becomes increasingly diffraction limited at higher frequencies, the DLFOV is likely to increase.  Since the D1 and D2 designs can significantly improve the available area in the FOV it is instructive to assess the shape of the mirrors that these designs require. The D1 mirrors are conic sections, an ellipsoid and a paraboloid, identical to those that define the CG, but the ellipsoid axis is tilted relative to the paraboloid axis. This tilt is chosen according to conditions outlined by Dragone\\cite{Dragone82}; in the case of the Archeops telescope the angle is 15$^{\\circ}$ (see Table~\\ref{Table-Archeops} and Figure~\\ref{Fig-D1-D2}). In the D2 design, localized corrections are applied to the shape of the mirrors of the D1 design. These corrections are designed to cancel coma near the center of the FOV\\cite{Dragone83}. The magnitude of the local surface corrections are given by $ K\\, r^{4}$, where $K$ is a constant that is different for each of the two mirrors and $r$ is the perpendicular distance from the the segment of the optical axis between the two mirrors, see Figure~\\ref{Fig-D1-D2}.  The constant $K$ depends on the distance between the mirrors along the optical axis, and for Gregorian telescopes the corrections are such that they curve the primary toward the secondary and the secondary away from the primary (see Figure~\\ref{Fig-D1-D2}). For the Archeops system that we have discussed in this paper $K = 3.54\\times10^{-9}$ and $7.76 \\times 10^{-8}$ cm$^{-3}$ for the primary and secondary, respectively. Given the sizes of the two Archeops mirrors the largest correction of the primary is 3.4 mm and is 2.4 mm for the secondary, values which are neither very large compared to the size of the mirrors, nor so small such as to make accurate machining difficult."
    },
    "0206/astro-ph0206027_arXiv.txt": {
        "abstract": "Intermediate-band photometry of the Hyades cluster on the {\\it Caby} system is presented for dwarf stars ranging from spectral type A through late K. A mean $hk$, $b-y$ relation is constructed using only single stars without anomalous atmospheres and compared to the field stars of the solar neighborhood. For the F dwarfs, the Hyades relation defines an approximate {\\it lower} bound in the two-color diagram, consistent with an [Fe/H] between +0.10 and +0.15.  These index-color diagrams follow the common convention of presenting stars with highest abundance at the bottom of the plot although the index values for the metal-rich stars are numerically larger. For field F dwarfs in the range [Fe/H] between +0.4 and --1.0, [Fe/H] = --5.6 $\\delta hk$ + 0.125, with no evidence for a color dependence in the slope. For the G and K dwarfs, the Hyades mean relation crosses the field star distribution in the two-color diagram, defining an approximate $upper$ bound for the local disk stars. Stars found above the Hyades stars fall in at least one of three categories: [Fe/H] below --0.7, [Fe/H] above that of the Hyades, or chromospherically active. It is concluded that, contrary to the predictions of model atmospheres, the $hk$ index for cool dwarfs at a given color hits a maximum value for stars below solar composition and, with increasing [Fe/H] above some critical value, declines. This trend is consistent, however, with the predictions from synthetic indices based upon much narrower Ca filters where the crossover is caused by the metallicity sensitivity of $b-y$. ",
        "introduction": "Among well-studied open clusters, the Hyades occupies a unique niche due to the combination of its proximity and its high metallicity \\citep{TAT97}. This rich sample of nearby stars of common age and composition provides a matchless testbed for any investigation of purely temperature-dependent trends at a given [Fe/H] among dwarfs with a wide range of mass. In particular, the high metallicity of the Hyades has led to its adoption as a reference point for zeroing and/or testing the metallicity scale for several photometric systems, including the {\\it UBV} \\citep{CA79}, the DDO \\citep{DEM77,TAT96}, and the Str\\\"omgren \\citep{CR75, SN89} systems. The purpose of this investigation is to add another, the {\\it Caby} system, to the long list of fundamental color relations defined by the Hyades, while investigating the effects of metallicity, and potentially of age, on the $hk$ index for a large sample of nearby field stars.  The {\\it Caby} system represents an extension of the traditional four-color, $uvby$ intermediate-band photometric system to a fifth filter centered on the H and K lines of Ca II. Details of the filter definition and design as well as the fundamental standards may be found in \\citet{ATT91}, while an extensive catalog of stars observed on the system and tied to the $b-y$ scale of \\citet{OLS93} may be found in \\citet{TAT95}. The filter was designed initially with metal-deficient stars in mind, as demonstrated by numerous applications to date on normal field stars \\citep{ATT91, AST92, ATT98, ATT00}, clusters \\citep{ATC95, RE00}, and variables \\citep{BD96, HI98}. Metallicity calibrations have been produced for both the metal-deficient giants \\citep{ATT98} and metal-deficient dwarfs \\citep{ATT00}, but preliminary analysis indicated that for dwarfs hotter than the sun, the $hk$ index, defined as {\\it (Ca-b)-(b-y)}, remains metallicity sensitive for stars of solar abundance or higher \\citep{TAT95}, a result consistent with the theoretical models of \\citet{SO93}. Because of the high metallicity of the Hyades relative to the typical star in the field of the solar neighborhood, it provides an ideal test of this prediction, as well as a means of probing the limits of its sensitivity.  A second issue of particular relevance for metal-rich stars is the role of chromospheric emission on the apparent strength of the $hk$ index. As one moves toward lower temperature, the expectation is that the absorption features will saturate, leaving an index which is almost exclusively dependent upon temperature. That the sensitivity of the {\\it hk} index to metallicity should decline near [Fe/H] = 0.0 for $(b-y)$ redder than 0.5 is readily apparent in the synthetic indices of \\citet{SO93}, though the models imply that even G and K dwarfs with abundances well above solar should exhibit larger $hk$ indices at a given color. However, line reversals triggered by chromospheric activity \\citep[see, e.g.,][]{VA80} may fill in the cores of the absorption features, making the star appear more metal-deficient than it actually is. Moreover, the strength of this line reversal has long been known to be age-dependent \\citep{WI63}. Since the Hyades is moderately young (less than 1 Gyr) compared to the average cool dwarf near the sun and contains some stars with an anomalous degree of chromospheric activity, it may provide some insight into this question. Finally, the majority of the stars known to be Hyades members have been studied in detail, thereby allowing us to look for additional photometric anomalies not tied to the metallicity of the stars.  Section 2 contains the details of the observations, their transformation to the standard $Caby$ system, and the merger of our $V$, $b-y$ data with an extensive array of published data on the Hyades dwarfs in an effort to minimize the potential effects of internal photometric scatter on the mean relations. In Sec. 3 we derive the Hyades mean relation in the {\\it hk, (b-y)} diagram and discuss the potential sources of intrinsic scatter in the two-color diagram. Sec. 4 compares the single-star Hyades relation to a large sample of nearby field stars, providing some insight into possible sources of the differences between expectation and reality, particularly at cooler temperatures. Sec. 5 contains a summary of our conclusions and suggestions for further work on the system in light of the Hyades anomalies. ",
        "conclusions": "An extensive sample of $Caby$ photometry of Hyades dwarfs from A through early M has been compiled and analyzed. For single stars and simple, composite, binary systems, the mean relation is well-defined over the color range from $b-y$ = 0.25 to 0.65. For the hotter portion of the color range, the effect of binarity on the two-color diagram is minimal. Among cooler dwarfs, the presence of a fainter, secondary star tends to shift the system above the mean relation, {\\it i.e.}, to lower $hk$ at a given $b-y$, simulating a lower [Fe/H], with a maximum offset between 0.05 and 0.10 mag. Stars with extreme degrees of chromospheric activity, particularly $BY Dra$ stars, appear anomalously metal-poor due to line-filling by emission within the Ca II H and K lines.  Among the F dwarfs, the Hyades relation defines a reliable lower bound to the distribution of field stars for a fixed [Fe/H] of +0.125. By comparison to the spectroscopic metallicity scale of \\citet{edv93}, it is found that the $hk$ index is linearly correlated with [Fe/H] over the range from +0.5 to --1.0, with almost twice the sensitivity to abundance changes compared to $m_1$. A simple test of the utility of $hk$ at super-metal-rich levels is provided by the sample of stars recently identified as having planets as compiled by the University of California Planet Search Team (http://exoplanets.org). Of the 76 systems listed, 16 are in the $hk$ Catalog \\citep{TAT95}. Of these, 6 are within the F-star color limits and have [Fe/H] between +0.17 and +0.34, with a mean of $+0.23 \\pm 0.08$. Line saturation does not appear to be a problem for the hotter stars.  In sharp contrast, as one extends the Hyades data toward cooler temperatures, the mean two-color relation crosses the distribution of nearby field stars, producing an approximate upper bound to the sample for stars in the color range from $b-y$ = 0.50 to 0.65. From an analysis of the stars that lie even higher than this relation in the $hk, b-y$ diagram, the primary source of the effect appears to be the high metallicity of the Hyades cluster. The implication is that the index does saturate at a given color for stars near solar abundance, but additional increases in [Fe/H] lower the index, placing super-metal-rich stars in the same region of the diagram as stars with [Fe/H] near --0.7 or lower.  An alternative explanation may be provided by the synthetic indices developed from model atmospheres by \\citet{SO93}. While the models imply that $hk$ should remain metallicity sensitive to [Fe/H] = +0.5 for all colors, a pattern almost identical to that found for $hk$ is generated for $C_{RV}$, the photospheric narrow-band index constructed from the data collected for the Mt. Wilson survey of stellar chromospheric emission \\citep{VA78}. \\citet{SO93} explain the crossover of $C_{RV}$ at cooler temperatures as primarily a result of the metallicity sensitivity of $b-y$, an effect that disappears with the use of $V-K$. Why this trend exists for an index based upon a filter over four times wider than those used in $C_{RV}$ remains a mystery, but it may be more of an indication of the model shortcomings than of a serious problem with the observed indices. In any case, the $hk$ index for cooler stars clearly succeeds in achieving the goal for which it was designed, the identification and calibration of stars of intermediate to extremely low [Fe/H]. This investigation adds two options to the list of uses, identification of stars with high levels of chromospheric activity, as expected, and the totally unexpected ability to isolate stars of Hyades abundance and higher. Three stars likely to occupy the last category and deserving of closer study are HD 57095, HD 76378, and HD 219495.  The question of the impact of modest, age-dependent variations in chromospheric emission remains unresolved, though indications from the very limited sample available to date are that the effect on $hk$ is small to negligible. Observations of the G and K dwarfs in nearby open clusters of comparable age but lower [Fe/H] than the Hyades could help to resolve this issue."
    },
    "0206/gr-qc0206060_arXiv.txt": {
        "abstract": "In a previous investigation, a model of three-body motion was developed which included the effects of gravitational radiation reaction. The aim was to describe the motion of a relativistic binary pulsar that is perturbed by a third mass and look for resonances between the binary and third mass orbits. Numerical integration of an equation of relative motion that approximates the binary gives evidence of such resonances. These $(m:n)$ resonances are defined for the present purposes by the resonance condition, $m\\omega=2n\\Omega$, where $m$ and $n$ are relatively prime integers and $\\omega$ and $\\Omega$ are the angular frequencies of the binary orbit and third mass orbit (around the center-of-mass of the binary), respectively. The resonance condition consequently fixes a value for the semimajor axis $a$ of the binary orbit for the duration of the resonance because of the Kepler relationship $\\omega=a^{-3/2}$. This paper outlines a method of averaging developed by Chicone, Mashhoon, and Retzloff which renders a nonlinear system that undergoes resonance capture into a mathematically amenable form. This method is applied to the present system and one arrives at an analytical solution that describes the average motion during resonance. Furthermore, prominent features of the full nonlinear system, such as the frequency of oscillation and antidamping, accord with their analytically derived formulae. \\\\ \\\\ \\textbf{Key words:} celestial mechanics, relativity, gravitational waves ",
        "introduction": "A model was established in a previous paper, `Gravitational Radiation Reaction and the Three Body Problem' (Wardell 2002), in which the motion of a binary system was studied. This binary system was subject to gravitational radiation damping and the gravitational influence of a third mass. The equations of motion for the relative motion of the binary were derived with certain approximations that would highlight the effects under investigation and also make analysis easier. For example, the motion of the three masses is taken to be planar and the center-of-mass of the binary moves in a fixed circular orbit around the third mass. Furthermore, the distance of the third mass from the binary's center-of-mass is taken to be substantially larger than the size of the relative orbit of the binary. The masses are considered to be point masses. After a scaling transformation that makes the variables dimensionless, one arrives at the following form of the equation of motion in Cartesian coordinates: \\begin{equation} \\frac{d^{2}r^{i}}{dt^{2}} = - \\frac{r^{i}}{r^{3}} -\\epsilon  K_{ij}r^{j}-\\delta R^{i}. \\end{equation} The variable $r^{i}$ represents the relative orbit, $K_{ij}(t)$ is the tidal matrix that retains the information about the tidal interaction between the third mass and the binary system, and $R^{i}$ is the radiation reaction term that expresses the radiation reaction force to desired order after iterative reduction. To avoid `runaway solutions' that can arise as the result of the radiation reaction perturbation which involves a fifth time derivative, one can apply the method of iterative reduction. This gives rise to a set of equations that is second order and resembles an equation of motion in Newtonian mechanics (Chicone et al. 2001). One recovers the differential equation for the appropriate Kepler problem if $\\epsilon=\\delta=0$.  Numerical analysis of the system has revealed that given appropriate initial conditions, one arrives at a result which shows resonance behavior in the relative orbit. When the graph of the Delaunay variable $L$, where $L=a^{1/2}$ such that $a$ is the semimajor axis of the osculating ellipse of the relative orbit, is plotted versus time one sees that the trend of semimajor axis decay can temporarily stop on average at a resonance. An oscillation occurs around a fixed average value for $L$ for the duration of this resonance. This resonance capture is indicative of an average balance of energy that leaves the binary system by way of gravitational waves and enters the system because of the tidal gravitational influence of the third mass.  This paper concerns itself with the behavior of the system when a resonance occurs; that is, when the resonance condition $m\\omega = n\\Omega'$ is satisfied, where $m$ and $n$ are relatively prime integers and $\\omega$ and $\\Omega'$ are the angular frequencies of the relative motion and tidal perturbation respectively. It turns out that the tidal perturbation frequency relates to the fixed third-body frequency as $\\Omega'=2 \\Omega$, where $\\Omega$ is the frequency of the fixed third-body motion.  An averaging method that was developed for the purposes of studying the effect of external incident gravitational waves on a binary system, can be applied to the equations of motion to analyze the behavior near a resonance (Chicone, Mashhoon \\& Retzloff 1997). The resultant averaged system of equations retains the `slow' variables which are left over after the system is averaged over a `fast' variable. The averaged set of nonlinear equations gives rise to a solvable system, whose solution approximates the actual solution for sufficiently small perturbation parameters over a certain time scale. ",
        "conclusions": ""
    },
    "0206/astro-ph0206094_arXiv.txt": {
        "abstract": "We report on $BVI$ CCD photometry of a field centered on the region of the intermediate-age open cluster NGC~2112 down to $V=21$. Due to the smaller field coverage, we are able to limit the effect of field star contamination which hampered in the past precise determinations of the cluster age and distance. This way, we provide updated estimates of NGC~2112 fundamental parameters. Having extended the photometry to the $I$ pass-band, we are able to construct  a colour-colour diagram, from which we  infer a reddening $E_{B-V}= 0.63\\pm0.14$ mag. The comparison of the Colour-Magnitude Diagram (CMD) with theoretical isochrones leads  to a distance of $850 \\pm 100$ pc, and an age of  $2.0 \\pm 0.3$ Gyr. While the distance is in agreement with previous determinations, the age turns out to be much better constrained and significantly lower than previous estimates. ",
        "introduction": "NGC~2112 (Collinder~76, C~0551-0031, OCL~509) is a northern open cluster of intermediate-age, located relatively far from the Galactic plane toward the anti-center direction ($\\alpha=05^{\\rm h}~53^{\\rm m}.9$, $\\delta=+00^{\\circ} 23^{\\prime}$, $l=205^{\\circ}.91$, $b=-12^{\\circ}.59$, J2000.). It is classified as a II2m open cluster by Trumpler (1930), and has a diameter of about $18^{\\prime}$, according to Lyng\\aa~ (1987). It is quite a poorly studied object, but rather interesting due  to its position in the disk and to its combination of suspected old age and low metal abundance, which would make it a noteworthy object to study in the framework of the chemical evolution of the Galactic disk (Carraro et al. 1998). Moreover it remained unstudied for long time due mainly, we guess, to the high contamination of field stars toward its direction which prevented precise estimates of its age and distance insofar (Richtler \\& Kaluzny 1989).\\\\ For these reasons, we decided to undertake a multicolour CCD study of the cluster, as presented in the present paper, which is the fourth of a series dedicated at improving the photometry of northern intermediate-age open clusters at Asiago Observatory. We already reported elsewhere on NGC~1245 (Carraro \\& Patat 1994), on NGC~7762 (Patat \\& Carraro 1995) and on NGC~2158 (Carraro et al.  2002).  \\noindent The plan of the paper is as follows. In Sect.~2 we summarize the previous studies on NGC~2112, while Sect.~3 illustrates the observation and reduction strategies. The analysis of the CMD is performed in Sect.~4, whereas Sect.~5 deals with the determination of cluster reddening, distance and age. Sect.~6 is dedicated to discuss the properties of the cluster in the context of the Galactic disk chemical evolution. Sect.~7 is devoted to analyze the geometrical structure and star counts and, finally, Sect.~8 summarizes our findings.  \\begin{figure} \\centerline{\\psfig{file=MC193fig1.eps,width=\\columnwidth}} \\caption{A DSS image of a region around NGC~2112 covered by the present study. North is up, east on the left.} \\label{mappa} \\end{figure} ",
        "conclusions": "We have presented a new CCD $BVI$ photometric study of the intermediate-age open cluster NGC~2112. The CMDs we derive are much cleaner than previous ones, and allow us to infer updated estimates of the cluster basic parameters. In detail, we find that:  \\begin{description} \\item $\\bullet$ the age of NGC~2112 is 2.0 Gyr, with a 15\\,\\% uncertainty; \\item $\\bullet$ the reddening $E_{B-V}$ turns out to be $0.63\\pm0.14$ mag; \\item $\\bullet$ we place the cluster at about 0.85 kpc from the Sun toward the anti-center direction; \\item $\\bullet$ we show that Brown et al. (1996) estimate of the metallicity is probably the most realistic one; \\item $\\bullet$ combining together age, distance and metallicity, we suggest that this cluster is a genuine member of the old thin disk population. \\end{description}   \\noindent As already noticed in the past, a proper motion study of NGC~2112 is really necessary to efficiently isolate cluster members from non-members and then derive robust estimates of the cluster chemical abundance."
    },
    "0206/astro-ph0206431_arXiv.txt": {
        "abstract": "Data obtained during five months of 2001 with the gravitational wave (GW) detectors EXPLORER and NAUTILUS were studied in correlation with the gamma ray burst data (GRB) obtained with the BeppoSAX satellite. During this period BeppoSAX was the only GRB satellite in operation, while EXPLORER and NAUTILUS were the only GW detectors in operation. No correlation between the GW data and the GRB bursts was found. The analysis, performed over 47 GRB's, excludes the presence of signals of amplitude $h\\ge 1.2 \\times 10^{-18}$, with 95\\% probability, if we allow a time delay between GW bursts and GRB within $\\pm400$  s, and $h \\ge 6.5 \\times 10^{-19}$, if the time delay is within $\\pm5$ s. The result is also provided in form of scaled likelihood for unbiased interpretation and easier use for further analysis. ",
        "introduction": "\\noindent One of the most important astrophysical phenomena still lacking an explanation is the origin of the celestial gamma-ray bursts (GRB). These are powerful flashes of gamma-rays lasting from less than one second to tens of seconds, with isotropic distribution in the sky. They  are observed above the terrestrial atmosphere with X-- gamma--ray detectors aboard satellites \\cite{fi,fis}. Thanks to the BeppoSAX satellite \\cite{Boella}, afterglow emission at lower wavelengths has been discovered \\cite{Costa,Jvp,Frail} and we now know that at least long ($>1 \\,$s) GRB's are at cosmological distances, with measured red shifts up to 4.5 (see, e.g., review by Djorgovski \\cite{Djorgovski} and references therein). Among the possible explanations of these events, which involve huge energy releases (up to $10^{54}$ erg, assuming isotropic emission), the most likely candidates are the collapse of a very massive star (hypernova) and the  coalescence of one compact binary system (see, e.g., reviews by Piran \\cite{Piran99} and M\\'esz\\'aros \\cite{Meszaros01} and references therein). In both cases emission of gravitational waves (GW) is expected to be associated with them (e.g. Ref. \\cite{ks}). According to several models, the duration of a GW burst is predicted to be of the order of a few milliseconds for a variety of sources, including the coalescing and merging black holes and/or neutron star binaries. Therefore GW bursts can be detected by the present resonant detectors, designed to detect GW through the excitation of the quadrupole modes of massive cylinders, resonating at frequencies near 1 kHz.  At the distances of the GRB sources ($\\approx 1$ Gpc), the GW burst associated with a total conversion of 1-2 solar masses should have amplitude of the order of $h  \\approx 3 \\times 10^{-22}$. The present sensitivity for 1 ms GW pulses of the best GW antennas with signal to noise ratio (SNR) equal to unity is $h \\approx 4 \\times 10^{-19}$ (see e.g. Ref. \\cite{piaamaldi}), which requires a total conversion of one million solar masses at 1 Gpc. However, although detection of a gravitational signal associated with a single GRB appears hopeless, detection of a signal associated with the sum of many events could be more realistic. Thus we launched a program devoted to studying the presence of correlations between GRB events detected with BeppoSAX and the output signals from gravitational antennas NAUTILUS and EXPLORER.  Searching for correlation between GRB and GW signals means dealing with the difference between the emission times for the two types of phenomena. Furthermore, there is also the fact to consider that the time difference can vary from burst to burst. In the present analysis we use an algorithm based on cross-correlating the outputs of two GW detectors (see \\cite{finn,arturo}), thus coping with the problem of the unknown possible time difference between GRB and GW bursts, and also of the unmodelled noise. ",
        "conclusions": "\\noindent Using for the first time a cross-correlation method applied to the data of two GW detectors, EXPLORER and NAUTILUS, new experimental upper limits have been determined for the burst intensity causing correlations of GW's with GRB's. Analyzing the data over 47 GRBs, we exclude the presence of signals of amplitude $h_{GW}\\ge 1.2 \\cdot 10^{-18}$, with 95\\% probability, with a time window of $\\pm~400~s$. With the time window of $\\pm~5$ s, we improve the previous GW upper limit to about $h=6.5~10^{-19}$.  The result is also given in terms of scaled likelihood and sensitivity bound, which we consider the most complete and unbiased way of providing the experimental information.  In a previous paper \\cite{vulcano} we had given more stringent upper limits, but this was under the hypothesis that the GW signals always occur at the same time with respect to the GRB arrival time. Here, instead, we only require that the time gap between the GRB and the GW burst be within a given time window. Similar comparison can be made with the AURIGA/BATSE result \\cite{cerdonio}, where an upper limit `` $h_{RMS} \\le 1.5~10^{-18}$ with C.L. 95\\% '' is estimated under the assumption that GW's arrive at the GRB time within a time window of $\\pm5$ s.  Finally, we remark that this method can be applied for any expected delay between GRB and GW, with appropriate time shifting of the integration window with respect to the GBR arrival time, according to the prediction of the chosen model."
    },
    "0206/astro-ph0206425_arXiv.txt": {
        "abstract": "There is growing  evidence that the majority of  the energy density of the universe is  not baryonic or dark matter, rather  it resides in an exotic  component   with  negative  pressure.   The   nature  of  this `quintessence' influences our view  of the universe, modifying angular diameter and luminosity distances. Here, we examine the influence of a quintessence component  upon gravitational lens time  delays.  As well as a static  quintessence component, an evolving equation  of state is also  considered.  It  is  found that  the  equation of  state of  the quintessence component  and its evolution  influence the value  of the Hubble's  constant derived  from gravitational  lenses.   However, the differences   between  evolving   and  non-evolving   cosmologies  are relatively small.  We undertake a  suite of Monte Carlo simulations to examine the potential constraints that  can be placed on the universal equation of  state from the  monitoring of gravitational  lens system, and demonstrate that  at least an order of  magnitude more lenses than currently known will have to  be discovered and analysed to accurately probe any quintessence component. ",
        "introduction": "The searches for supernovae at cosmological distances have proved very successful,  providing evidence  that, while  topologically  flat, the majority  of energy  in  the Universe  is  in the  form  of an  exotic component with  negative pressure (Riess  et al.  1999;  Perlmutter et al.   1999).  The  recent  identification of  a  supernova at  $z=1.7$ (Riess et  al.  2001)~\\footnote{It should be noted,  however, that the influence  of gravitational  lensing  on SN1997ff  needs  to be  fully addressed before  its true cosmological significance  can be addressed (Lewis  \\&  Ibata 2001;  Moertsell,  Gunnarsson  \\& Goobar~2001)}  has provided further weight to these  claims (Turner \\& Riess 2001), which suggest that this component may differ from the classical cosmological constant $\\Lambda$.  Termed `quintessence', or more colloquially `dark energy', this has an equation of state of the form $P = w \\rho$, where $P$  is the pressure  and $\\rho$  the density.   A $w  < -\\frac{1}{3}$ opposes the action of gravity and drives the cosmological expansion to accelerate.  Linder (1988a; 1988b) has examined the physical nature of various    quintessence   components;    with   $w=0$    equating   to non-relativistic  matter (dust),  $w=\\frac{1}{3}$ being  radiation and $w=-1$,  a classical  cosmological constant.   More  exotic components are;   massless   scalar   fields   $w=1$,  cosmic   string   networks $w=-\\frac{1}{3}$,    and     two-dimensional    topological    defects $w=-\\frac{2}{3}$.    As  well   as  the   supernova   programs,  other approaches,  such  as  gravitational  lensing  statistics  (Cooray  \\& Huterer  1999), geometrical  probes  of the  ${\\rm Ly_\\alpha}$  forest (Hui, Stebbins  \\& Burles 1999) and galaxy  distributions (Yamamoto \\& Nishioka  2001), and  classical angular-size  redshift tests  (Lima \\& Alcaniz  2001), will  provide  complementary probes  of the  universal equation of state.  The value of  the quintessence component, $w$, influences  our view of the  universe, modifying  the various  distances used  in  mapping the cosmos.   This paper  concerns itself  with  the influence  of $w$  on angular   diameter   distances,   especially   in  relation   to   the determination of  the Hubble's constant  from the measurement  of time delays in gravitational lens  systems.  Unlike local determinations of Hubble's constant (e.g. Freedman et al. 2001), the cosmological nature of  gravitational lenses  means that  they are  more sensitive  to the underlying cosmological  parameters.  Section~\\ref{background} briefly covers the  basic formulae for generalized  angular diameter distances in  quintessence  cosmologies,  while  in  Section~\\ref{timedelay}  we consider  the influence  of $w$  on  the determination  of $H_o$  from lensed systems.  Section~\\ref{evolve}  extends this analysis to simple models     of    an     evolving    quintessence     component.     In Section~\\ref{montecarlo}  a  series  of  Monte Carlo  simulations  are undertaken to  estimate the efficacy  of this approach in  probing the cosmological  equation of state,  while in  Section~\\ref{speculate} we speculate   on   the   possability   that  current   observations   of gravitational lens  systems may suggest that  $w<-1$.  The conclusions of this study are presented in Section~\\ref{conclusions}.  \\begin{figure*} \\centerline{ \\psfig{figure=fig1.ps,height=14.5cm,angle=270.0} } \\caption{The relative  time delays for  several flat cosmologies  as a function of  the equation  of state,  $w$.  Each curve  is is  for the lowest value of  the time delay (the redshift  of the lens-source pair is given in brackets).  Each panel presents a different combination of $\\Omega_m$ and  $\\Omega_w$; note that $w=0$ corresponds  to a universe composed entirely of matter in all cases.} \\label{fig1} \\end{figure*} ",
        "conclusions": "This paper  has investigated the  role of a quintessence  component on angular  diameter  distances,  specifically  their  influence  on  the determination  of Hubble's  constant from  the measurement  of  a time delay in multiply imaged quasars.  For  flat universes,  with an  unevolving quintessence  component, its seen that, for a gravitational lens system in which the time delay has been measured,  the resultant Hubble's constant is  dependent upon the value  of the  equation of  state parameter  $w$.   Interestingly, the dependence  of the  determined value  of  the Hubble's  constant as  a function of $w$  possesses a minimum which is  independent on the lens and source redshift.  Several models of evolving quintessence were also examined, consisting of  a linear evolution  of the  equation of  state with  redshift. The cosmologies resulted  in significantly different forms  of the angular diameter distance. Hence, our view of the cosmos would be different in the various cosmologies. When  considering the specific combination of angular   diameter   distances   that  constitute   the   cosmological contribution  to the gravitational  lensing determination  of Hubble's constant, it is seen that the resulting variations between cosmologies is  very  small, a  matter  of  only a  few  percent,  relative to  an unevolving case with the same present day constitution.  A number of  Monte Carlo simulations of the  determination of Hubble's constant and the quintessence  equation of state, $w$, were undertaken to explore  the efficacy  of this approach.   These revealed  that the present situation with only a handful of lensed systems does not allow an accurate determination of the cosmic equation of state, and that at least an order of magnitude  more lenses are truly required to provide a reasonably  robust determination  of the underlying  cosmology.  The next generation of all-sky surveys are presently underway (e.g.  Sloan Digital Sky  Survey) or  are being planned  (e.g.  VISTA,  PRIME), and these  datasets will  greatly increase  the number  of  lensed quasars available for  monitoring studies.  Cooray \\&  Huterer (1999) estimate that  $\\sim 2000$  lensed quasars  will  be identified  from the  SDSS database  alone, and a  much larger  number can  be expected  from the deeper VISTA and  PRIME surveys. The number of  these sources amenable for follow-up monitoring campaigns will naturally be much smaller, but one  can confidently  expect a  sample of  several hundred  systems to eventually become  available.  However,  given the effort  required to first find such systems, as well as monitor them to determine the time delays and the modeling procedure, it is likely that $w$ will be first determined using  one of the other various  techniques currently being proposed.  We conclude, therefore, that gravitational lens time delays are likely to prove poor probes of the universal equation of state.  \\newfont{\\afont}{cmfi10}"
    },
    "0206/astro-ph0206339_arXiv.txt": {
        "abstract": "Weak lensing of high-redshift Type Ia supernovae induces an external dispersion in their observed standard candle brightnesses, comparable in magnitude to the intrinsic dispersion for redshifts $z>1$.  The same matter fluctuations responsible for the magnification of distant supernovae also generate shear in the images of background galaxies.  We investigate the possibility of using lensing shear maps constructed from galaxies surrounding the supernovae as a means of correcting the lensing-induced magnification dispersion.  We find that a considerable fraction of the lensing dispersion derives from sub-arcminute scales, which are not probed by shear maps smoothed on arcminute scales.  We thus find that weak lensing shear maps will be of only limited value in reducing the weak lensing magnification fluctuations of supernovae. ",
        "introduction": "High redshift Type Ia supernovae provide an excellent means of studying the expansion history of the universe~\\citep{hzss,scp}. It is estimated that the intrinsic dispersion in supernova luminosities can be calibrated to $\\approx0.15$ mag, and perhaps in the future to $0.1$ mag~\\citep{snap}, making them excellent standard candles.  For supernovae at redshifts $z<1$, this intrinsic dispersion sets the limiting accuracy with which supernovae may be used to measure distances.  For higher redshifts ($z\\ga1$), however, gravitational lensing by random fluctuations in the intervening matter distribution induces a dispersion in supernova brightness comparable to the intrinsic dispersion \\citep{frieman97,hw98,holz98,wang}, degrading their value as standard candles. These magnification fluctuations have zero mean, and so may be averaged away with sufficient numbers of supernovae \\citep{wangbin}.  However, the additional dispersion means that more supernovae are required than for low redshift samples to achieve a given signal to noise.  It would be of great utility to determine the gravitational lensing magnification of each individual supernova. This would allow a correction of the observed brightnesses of the supernovae, and therefore improve their use as standard candles. Such a correction would be equivalent to obtaining a larger sample of supernovae, for free. In addition, measuring the gravitational lensing distribution at high redshift can be an important probe of the dark matter~\\citep{ms99,seljakholz,ben99,holz01}.  One means of achieving this would be an inspection of the foreground galaxies for each supernova.  For example, SN1997ff at $z=1.7$ has several foreground galaxies in its vicinity, leading to a magnification possibly as large as 0.4 magnitudes~\\citep{ibata,sn1997ff,moertsell}.  If the magnification factor could be accurately estimated from the foreground galaxy images, then the supernova brightness could be corrected to its unlensed value.  The correction factor depends strongly on uncertain properties of the galaxies' mass distributions (illustrated by the controversy over the extent of lensing of SN1997ff), and would miss possibly important contributions from dark halos.  Furthermore, since such corrections would primarily shift highly magnified SNe to lower brightnesses, while leaving demagnified SNe unaffected, it would bias the resultant Hubble diagram. It is apparent that direct identification of individual lenses does not robustly determine the lensing magnification.  It is also possible to correlate, in a statistical manner, the foreground galaxy number density close to the lines of sight to supernovae with the lensing effects on these supernovae~\\citep{ben01}, but these statistical results do not help us ``correct'' any given individual supernova.  An alternative method for correcting lensing magnification is to utilize weak lensing maps constructed from shear measurements of background galaxies. The same matter fluctuations responsible for the magnification of supernovae also lead to shearing of galaxy images.  High redshift SNe are discovered by repeated exposures of wide fields, which when co-added provide extremely deep images of the galaxies surrounding the supernovae.  Such deep, wide field images are well-suited for measurement of weak lensing shear.  It is thus natural to hope that mass reconstruction from shear measurements of the surrounding fields might allow for the correction of weak lensing magnification, restoring the supernovae to their intrinsic brightnesses.  A perfect measurement of the shear field at the redshift of a given supernova would allow for a perfect reconstruction of the projected mass surface density (modulo the mass-sheet degeneracy, which should be unimportant for large enough fields). From this mass surface density it is possible to calculate the lensing magnification, and therefore perfectly account for (and correct) the lensing effects on the observed brightness. Perfect shear maps are unavailable, however, and therefore our ability to infer the magnification is compromised.  In this paper, we investigate how well weak lensing reconstruction can correct the brightnesses of distant supernovae.  The basic scheme is as follows. A supernova occurs in a given field, and its peak apparent magnitude is observed and calibrated, using some variant of the \\citet{phillips} relation. Then the (co-added) field containing the supernova is used to estimate the local shear at the supernova's location by averaging over a smoothing angle $\\theta$. The shear map is then converted to an effective convergence map using some reconstruction algorithm such as that of \\citet{ks}, and the derived convergence is used to correct the supernova's standard candle brightness. In the following section we estimate the variance in convergence for point sources given knowledge of the smoothed shear map, $\\langle\\kappa^2\\rangle_\\gamma$, which is a direct measure of the improvement such an approach can offer. We find that useful corrections require very large background source galaxy densities, and that this method is therefore of only marginal utility. ",
        "conclusions": "From Figure~\\ref{rsq} it is apparent that shear maps will be of limited value in reducing the lensing dispersion of supernova brightnesses, unless the number density of background galaxies is great enough to permit sufficiently small smoothing angles.  For example, if the number density of background galaxies is as high as ${\\bar n}=10^6/{\\rm deg}^2$, then at best we find a value for the cross-correlation coefficient of $r^2\\simeq0.35$.  If the uncorrected convergence variance is $\\langle\\kappa^2\\rangle=0.0036$, then we find the variance for the corrected supernovae convergence to be $\\langle\\kappa^2\\rangle_\\gamma=(1-r^2)\\langle\\kappa^2\\rangle=0.0023$. This yields an rms magnification of $0.1$, which is a 20\\% improvement over the uncorrected value of $0.12$.  This represents the optimal case discussed above, where all of the source galaxies are at the same redshift as that of the supernova. Including the expected spread in galaxy redshifts provides a more reasonable estimate of $r^2\\simeq 0.2$, giving a reduction in the rms magnification of the supernova of around $10\\%$.  It is to be emphasized that this is an improvement (reduction) in the width of the observed supernova magnification distribution, and not a change in the mean (which remains at $\\mu=1$). In addition, the {\\em intrinsic}\\/ dispersion in supernova luminosities causes a further contribution to the observed rms standard candle magnification luminosity.  Tomographic information can do little to ameliorate the situation.  The simplest approach would be to confine the shear analysis to source galaxies in a slab in redshift space centered on the supernova. By doing this one moves up and to the left of the solid curve in Figure~\\ref{rsq}, trading off increased shot noise for more effective lensing information. It is apparent from the Figure that there is a net improvement if $\\ga10\\%$ of the galaxies are at similar redshifts to that of the supernova.  More inspired schemes might attempt to employ the information contained in galaxies at all redshifts; regardless, for a given effective galaxy density the theoretical limit is still bounded by the dashed curve in Figure~\\ref{rsq}. Our model for the distribution of galaxy source images in redshift space is particularly simple---more realistic models (e.g. with dependence on survey depth) may also push one closer to the dashed curve.  The estimate presented here is optimistic in that we assume that the smoothed convergence field may be directly measured.  In reality, the shear field is measured, and then converted to a convergence map~\\citep{ks}. Even in this optimistic approximation, at best meager returns are expected from the construction of shear maps of surrounding galaxies. An additional caveat is that we have assumed the noise is dominated by Poisson noise in the number of source galaxies. Additional systematic errors, such as imprecise measurement of the point spread function, only worsen the decorrelation. Note that our conclusions are sensitive to the shape of the convergence angular power spectrum. If in reality the power spectrum is unlike that of Figure~\\ref{angpow}, and instead has far less small-scale power, then galaxy shear may turn out to be a much more powerful tool for correcting weak lensing of supernovae.  At present, it appears that there is significant small-scale power \\citep{dk1,dk2}, consistent with the values assumed here.  Future wide field surveys like the LSST\\footnote{http://www.lssto.org} or SNAP\\footnote{http://snap.lbl.gov} will directly measure the convergence angular power on some of the relevant scales, so it will be possible to check whether the assumptions made here are valid.  Assuming, however, that the power spectrum does not significantly depart from that which we have used, the prospects for using galaxy shear to correct supernova brightnesses appear bleak. Given the danger of introducing unknown biases in the resulting distance-magnitude relation, it is unclear whether future supernova surveys should attempt the use of lensing shear maps to correct for magnification of supernova brightnesses."
    },
    "0206/astro-ph0206080_arXiv.txt": {
        "abstract": "The RS CVn-type binary $\\sigma$ Geminorum was observed during a large, long-duration flare simultaneously with {\\it XMM-Newton} and the VLA. The light curves show a characteristic time dependence that is compatible with the Neupert effect observed in solar flares: The time derivative of the X-ray light curve resembles the radio light curve. This observation can be interpreted in terms of a standard flare scenario in which accelerated coronal electrons reach the chromosphere where they heat the cool plasma and induce chromospheric evaporation. Such a scenario can only hold if the amount of energy in the fast electrons is sufficient to explain the X-ray radiative losses. We present a plausibility analysis that supports the  chromospheric evaporation model. ",
        "introduction": "There is compelling evidence that high-energy processes and high-energy particles play a pivotal role in the energy release, energy transport, and plasma heating during solar flares (see review by \\citealt{hudson95}). A standard scenario proposes that electrons (perhaps also ions) are accelerated in the corona in the course of magnetic reconnection. As the electrons travel along closed magnetic fields, those  with large pitch angles and sufficient energy  (typically several 100~keV) lose a small part of their energy as  gyrosynchrotron radio emission. The bulk kinetic energy of the accelerated electrons, however, is carried to the chromosphere where it is deposited by electron-ion collisions. The collision of the beam with the dense plasma reveals itself by non-thermal hard X-ray radiation (HXR, typically between 10$-$100~keV) that, however, constitutes only a small fraction ($\\sim 10^{-5}$)  of the total energy loss. The bulk energy is transformed into heat, producing an overpressure in the chromosphere as the gas cannot radiate away the energy influx sufficiently rapidly. As a consequence, the gas evaporates explosively into the corona as a $\\sim 10^7$~K plasma visible in X-rays during the gradual phase of a solar flare \\citep{dennis88}. The observed gyrosynchrotron radio luminosity $L_{\\mathrm{R}}$ and the hard X-ray luminosity $L_{\\mathrm{HXR}}$ are, to first order, proportional to the instantaneous number of fast electrons and therefore to the power $\\dot{E}$ injected into the system, while the slowly variable soft X-ray luminosity $L_{\\mathrm{X}}$ is roughly proportional  to the accumulated total energy $E$ in the hot coronal flare plasma. One therefore expects that \\begin{equation}\\label{neupert} {d\\over dt} L_{\\mathrm{X}}(t) \\propto L_{\\mathrm{HXR}}(t) \\propto L_{\\mathrm{R}}(t) , \\end{equation} a relation that is commonly known as the `Neupert Effect' \\citep{neupert68,dennis93}. Although significant deviations from this scenario have been observed in solar flares (e.g., heating starting before any hard X-rays can be detected, or absence of one of the emission types discussed above), there is strong support for several features of this model in the majority of solar flares \\citep{dennis93}. For example, the coincidence to within a fraction of a second of HXR brightenings at a pair of magnetic loop footpoints that are separated by $\\sim 10^9$~cm  requires non-thermal particle velocities \\citep{sakao94}.  It is thought that the same mechanisms should operate in stellar flares, although flares on some classes of stars deviate considerably from the proposed solar analogy. In particular, giant flares on RS CVn-type binaries may require mechanisms unknown on the Sun. RS CVn binaries, commonly consisting of a giant or subgiant primary with a main-sequence or subgiant companion in a close orbit, are sources of luminous radio and X-ray emission \\citep*{drake89}. Very long flare time scales and radio source sizes of order of the intrabinary distance have been modeled in terms of giant dipole-like magnetospheric structures \\citep*{morris90, jones94} into which high-energy particles are injected from a flare site \\citep{mutel85}, and where they lose most of their energy by radiation. Coordinated observations in X-rays and radio are required to study the importance of high-energy electrons in the heating mechanism. \\citet{hawley95} reported a Neupert effect-like behavior during a large flare on the dMe star AD Leo, where optical (U band) and EUV emissions were used as proxies for the radiation from high-energy electrons (radio, HXR) and from the thermal plasma,  respectively. \\citet{guedel96} discussed the first stellar Neupert effect seen in the  radio and X-ray bands on the dM5.5e binary UV Cet, finding similar timing and similar energy budgets as in solar gradual (``type C'') events. ",
        "conclusions": "We note that, within the framework of our simplified model assumptions,  our energy estimate is conservative. The lifetime $\\tau$ of the non-thermal electron population has been adopted at the highest possible value compatible with the radio light curve (i.e., approximately equal to the shortest decay time scale in the observed radio emission). Equation~\\ref{integral} shows that a shorter possible lifetime requires a proportionately larger total energy input. In that case, the light curve variability is controlled by the time scales of the particle accelerator. If electrons get lost after one magnetic-loop crossing time, then $\\tau \\approx$ a few seconds for typical loop sizes, and the rapid replenishment of electrons requires an energy input into the system up to 3 orders of magnitude higher than estimated above. We conclude that not only does the relative timing between radio and X-ray emissions support the chromospheric evaporation scenario, but the total energy content in the injected high-energy electrons could easily satisfy or largely exceed the requirements set by the observed X-ray losses.  The observation described here provides strong support for the chromospheric evaporation scenario in a star that may maintain a corona considerably different from the Sun's (e.g., containing much larger magnetic loops, confining much hotter and perhaps also denser thermal plasma, magnetic field lines that may be arranged in the form of large global dipoles, possibly also between the companion stars, etc). In retrospect, we find a similar timing between radio and X-ray flare events in some previously published light curves, although the Neupert effect was not discussed. Most evidently, radio emission peaking before the soft X-rays, thus suggesting the presence of a Neupert effect,  can be seen in the examples presented by \\citet{vilhu88}, \\citet{stern92},  \\citet{brown98}, and \\citet{ayres01}. Clearly, differing behavior has been noted as well. First, the Sun shows the Neupert effect most reliably in the class of impulsive flares, whereas 50\\% of all gradual flares, often related to energy release at high coronal altitudes, show a different behavior \\citep{dennis93}. It is possible that the connectivity of magnetic fields between the high corona and the chromospheric regions is different in these cases, impeding the free flow of electrons and consequent chromospheric evaporation. Thermalization of fast electrons could also occur in the corona already if the travel distances are long enough and the densities high enough. Stellar counter-examples of the Neupert effect include an impulsive optical flare with following gradual radio emission \\citep{vdoord96}, gyrosynchrotron emission that peaks after the soft X-rays \\citep{osten00}, and an X-ray depression during strong radio flaring \\citep{guedel98}. Note also that complete absence of correlated flaring has been observed at radio and UV wavelengths (e.g., \\citealt{lang88}).   Evidence for  chromospheric evaporation in an RS CVn binary system is potentially important to understand to what degree the solar analogy can be applied to such stellar systems. Although magnetospheric sizes as measured in radio waves are several times larger than the Sun, VLBI observations have suggested that flares start out in compact, unresolved cores that are well localized in magnetic active regions close to the stellar surface \\citep{mutel85}. Our observations of a Neupert effect strongly suggest solar analogy in the physics of energy release and transport in this binary system, at least for the large flare reported here."
    },
    "0206/astro-ph0206049_arXiv.txt": {
        "abstract": "We present radio observations of 19 candidate compact steep-spectrum (CSS) objects selected from a well-defined, complete sample of 52 B2 radio sources of intermediate strength. These observations were made with the VLA A-array at 4.835 GHz. The radio structures of the entire sample are summarised and the brightness asymmetries within the compact sources are compared with those of the more extended ones, as well as with those in the 3CRR sample and the CSSs from the B3-VLA sample. About 25 per cent of the CSS sources exhibit large brightness asymmetries, with a flux density ratio for the  opposing lobes of $>$5, possibly due to interaction of the jets with infalling material. The corresponding percentage for the larger-sized objects is only about 5 per cent.  We also investigate possible dependence of the  flux density asymmetry of the lobes on redshift, since this might be affected by more interactions and mergers in the past. No such dependence is found.  A few individual objects of interest are discussed in the paper. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206321.txt": {
        "abstract": "An atlas of far-ultraviolet spectra of 45 Galactic OB stars observed with the {\\it Far Ultraviolet Spectroscopic Explorer} is presented. The atlas covers the wavelength region between 912 and 1185~{\\AA} with an effective spectral resolution of 0.12~{\\AA}. Systematic trends in the morphology and strength of stellar features are discussed. Particular attention is drawn to the variations of the {\\ciii~$\\lambda$1176}, {\\siv~$\\lambda\\lambda$1063, 1073}, and {\\pv~$\\lambda\\lambda$1118, 1128} line profiles as a function of temperature and luminosity class; and the lack of a luminosity dependence associated with {\\ovi~$\\lambda\\lambda$1032, 1038}. Numerous interstellar lines are also identified. ",
        "introduction": "The far-ultraviolet (FUV; 900--1200~{\\AA}) region of the spectrum contains an enormous number of spectral features attributable to resonance lines and transitions between excited states. These transitions are due to a variety of atomic and molecular species, which include elements that are cosmically abundant and some that are comparatively rare. Collectively, these lines diagnose a wide range of ionization and excitation conditions, and can therefore provide extremely detailed information about the physical conditions that exist in astrophysically interesting environments like stellar atmospheres, the interstellar medium (ISM), and the intergalactic medium (IGM).  Unfortunately, observations in the FUV are also very challenging from a technical standpoint, both because of the low reflectivity of optical surfaces at such short wavelengths and the need for high spectral resolution to minimize confusion from line blending. Consequently, despite the many scientific incentives, the FUV window has been underutilized. Before 1999, the main forays into the FUV were limited to the {\\it Copernicus} satellite {\\citep[1972--1981;][]{rog73}} and a series of comparatively short duration shuttle-based missions: the Hopkins Ultraviolet Telescope {\\citep[HUT;][]{dav92}} which flew on the {\\it Astro-1} and {\\it Astro-2} missions; and the Interstellar Medium Absorption Profile Spectrometer (IMAPS; Jenkins, Reale, \\& Zucchino 1996), the Berkeley Extreme and Far Ultraviolet Spectrometer {\\citep[BEFS;][]{hur98}}, and the T{\\\"u}bingen Echelle Spectrograph {\\citep[TUES;][]{barn99}}, which flew on the {\\it ORFEUS-SPAS I} and {\\it II} missions.  The launch of the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\fuse}) in 1999 June rectified this situation. Since then, {\\fuse} has provided routine access to the entire FUV waveband with high spectral resolution and exceptional sensitivity, which in turn has permitted observations of a substantially larger target pool than was available to {\\it Copernicus} (due to sensitivity limitations) or the shuttle-based missions (due to time constraints). In particular, spectra of more than 200 Galactic OB-type stars covering most spectral types\\footnote[10]{See \\citet{wal02} for a description of the criteria that define the new O2 spectral class.} and luminosity classes between O2 and B9 have been obtained as part of the various programs implemented by the Principal Investigator (PI) Team.  We have selected a subset of these objects for presentation in a high-resolution FUV spectral atlas. The principal aim of this atlas is to show the general behavior of the most prominent stellar lines as a function of temperature and luminosity class, and also to illustrate the rich diversity of the interstellar spectrum. This work has several immediate applications, such as line identification, the characterization of hot-star winds, FUV extinction, population synthesis, and the interpretation of young stellar populations in distant galaxies. More fundamentally, we hope that this atlas will help familiarize researchers with the FUV region of the spectrum. It complements existing atlases based on spectra obtained with {\\it Copernicus\\/} {\\citep[e.g.][]{snow76, snow77, wal96}} both by including the spectral region between 910 and 1000\\AA\\ and by enlarging the sample of objects earlier than B3, and serves as a companion to the {\\fuse} atlas of OB-type stars in the Magellanic Clouds \\citet{mcatlas}.  The remainder of this paper is organized as follows. The observational material and data processing are summarized in {\\S 2}, while {\\S 3} provides general comments on the organization of the atlas. Section~4 gives a basic overview of the principal interstellar lines found in the FUV, followed by {\\S 5} and {\\S 6}, which describe the major trends exhibited by the stellar features as a function of temperature and luminosity class, respectively. Concluding remarks are given in {\\S 7}. ",
        "conclusions": "The primary result emerging from this atlas is a refined understanding of the trends in the strength and morphology of the wide variety of spectral lines found in FUV spectra of OB-type stars. It is remarkable that these FUV features -- the most prominent of which arise in the stellar wind -- vary so smoothly according to spectral type and luminosity class, which are defined by photospheric lines in the optical region of the spectrum. These smooth variations were already observed from IUE data by Walborn, Nichols-Bohlin, \\& Panek (1985). An obstacle to the detection of these correlations is the accuracy of the optical spectral classifications. Some of the classifications for the Galactic targets accessible to {\\fuse} are quite dated, and may not have been classified according to the same criteria currently used for brighter standards. To correct this situation, we are currently collecting new optical spectra for some of these stars, so that their published spectral types can be verified.  This atlas also confirms previous results {\\citep[e.g.][]{snow76, wal96}} concerning the sensitivity of various FUV lines to physical conditions in the stellar atmosphere. In particular, the strength and morphology of the {\\ciii~$\\lambda$1175.6} line is shown to be a strong function of both effective temperature and luminosity. Similar behavior is demonstrated for the {\\siv~$\\lambda\\lambda$1063, 1073} and {\\pv~$\\lambda\\lambda$1118, 1128} lines.  The {\\ion{O}{6}~$\\lambda\\lambda$1032, 1038} wind profiles, which persist to early B-type stars, exhibit few variations as a function of luminosity.  Finally, this spectral atlas illustrates directly the treasure-trove of astrophysical information that lies encoded in FUV spectra. The broad wavelength coverage, sensitivity, and good spectral resolution of {\\fuse} provide many opportunities to study the atmospheres of early-type stars and the intervening ISM in fundamentally new ways. As one application of this new capability, we are currently adding {\\fuse} spectra of Galactic OB-type stars to the databases used by population synthesis codes {\\citep[e.g., Dionne \\& Robert, in preparation;][]{ star99}} in order to use the FUV line sensitivities noted above to refine studies of young stellar populations in distant galaxies."
    },
    "0206/astro-ph0206033_arXiv.txt": {
        "abstract": "Star-forming regions have been the targets of X-ray observations since the dawn of satellite X-ray astronomy.  The increase in sensitivity and/or spatial resolution offered by \\XMM\\ and \\Ch\\ allows a dramatic improvement, both qualitative and quantitative, on our knowledge of high-energy phenomena in these regions and the underlying physical processes.  We summarize here some recent developments:  the Orion Nebula Cluster and its 1000+ stellar X-ray sources; Herbig-Haro objects and their high-speed shocks; protostars, brown dwarfs and their unusual magnetic activity; and the discovery of diffuse X-ray emission from HII regions, presumably related to strong winds from massive stars.  The role that future X-ray missions may play in the field is already starting to be visible. ",
        "introduction": "Star-forming regions have been known to be associated with X-ray emission for over 30 years, beginning with the discovery of an extended source coincident with the Orion Nebula (M42) using the \\U\\, {\\sl ANS} and {\\sl SAS-3} satellites (Den Boggende et al.  1978, Bradt \\& Kelley 1979). Possibilities to produce X-rays from star forming regions include magnetic activity from lower mass pre-main sequence stars; thermalization of the high velocity winds of higher mass OB stars, either close to the star or where at a wind termination shock; and supernova remnants from past generations of OB stars (Figure \\ref{fig:sfr_diagram}). It took the advent of the first {\\it imaging} X-ray satellite, the \\E\\ Observatory, to realize that stars were the ``true'' X-ray emitters in the Orion Nebula. Progress was rapid from the start, owing to the large field-of-view, 1 sq.  deg.  or more for satellites such as \\E\\ and \\R, allowing to detect dozens of stellar sources in a single exposure of nearby star forming regions (see review by Feigelson \\& Montmerle 1999, henceforth FM).  \\begin{figure}[ht] \\begin{center} \\epsfig{file=tmontmerle-E2_fig1.eps, width=8.5cm} \\end{center} \\caption{Diagram of the expected X-ray components from a giant molecular cloud with a blister H{\\sc II} region, embedded young star clusters and distributed star formation. Symbols: {\\Large\\bf $\\star$} = OB stars; {\\bf $\\times$} = Herbig Ae/Be stars; $\\bullet$ = T Tauri stars; {\\bf +} = protostars; {\\it squares} = X-ray binary system.  The hatched region outside of the cloud represents a supernova remnant, and shaded regions within the cloud represent partially ionized X-ray dissociation regions (see Feigelson 2001). } \\label{fig:sfr_diagram} \\end{figure}    In parallel with the development of X-ray satellites, major progress was being achieved in sensitive solid-state detectors for ground-based telescopes in the IR and mm ranges.  This led to the discovery of {\\it circumstellar material}, accretion disks and envelopes around young stars and protostars (see, e.g., Andr\\'e \\& Montmerle 1994).  High-angular resolution imaging, by, e.g., the \\HST\\ and mm interferometers, showed that young stars and protostars (``Young Stellar Objects'': YSOs) lose mass in the form of jets and outflows, collimated perpendicular to the circumstellar disks.  Thus accretion and ejection are currently viewed as correlated phenomena, somewhat paradoxically required to build up a star. Fig.~\\ref{fig:HH30} gives a spectacular example, obtained with the \\HST, of a disk-jet system associated with a very young star.   \\begin{figure}[ht] \\begin{center} \\epsfig{file=tmontmerle-E2_fig2.eps, height=7cm} \\end{center} \\caption{HST WFC2 optical view of the YSO HH30 in Taurus in 2000. The circumstellar, flared disk, of radius $\\sim 200$ AU, is seen almost edge-on. A bright, collimated jet is expelled from the central regions.} \\label{fig:HH30} \\end{figure}   It is widely believed that magnetic fields are required to explain, and perhaps even cause, the ejection of material, from the disk, and/or from the central, growing star (e.g., K\\\"onigl \\& Pudritz 2000). Fig.~\\ref{fig:MagConf} illustrates various star-disk magnetic configurations from the literature. Recent calculations indicate that YSO X-rays which, as on the Sun, are produced in violent magnetic reconnection events, may be the principal ionization source at the base of these jets (Shang et al. 2002).  We thus see that in YSOs X-rays, circumstellar matter, and magnetic fields are somehow intimately interrelated.   \\begin{figure}[!ht] \\begin{center} \\epsfig{file=tmontmerle-E2_fig3.eps, width=8.5cm} \\end{center} \\caption{Various star-disk magnetic configurations, taken from the literature. Top left: Shu et al. (1997). Top right: Hirose et al. (1997). Bottom left: Hartmann (2000). Bottom right: Ferreira et al. (2000).} \\label{fig:MagConf} \\end{figure} ",
        "conclusions": "As in other fields of X-ray astronomy, \\Ch\\ and \\XMM\\ have opened a new era for the study of star-forming regions and the early stages of stellar evolution.  This era is only starting, since at the time of this Conference only a limited amount of results are available. However, the selection we have presented above already demonstrate the new possibilities:  $\\bullet$ In dense young clusters, like the ONC, several hundreds of low-mass stars can be detected and identified simultaneously in one image, from massive OB stars down to substellar objects which will evolve into brown dwarfs;  $\\bullet$ Protostars are detected with a high efficiency. Most detected protostars are Class I (evolved, age $\\sim 10^5$ yrs), along with two candidate Class 0 (young, age $\\sim 10^4$ yrs);  $\\bullet$ X-ray flares testify to the intense magnetic activity (stellar surface or star-disk reconnection events) of YSOs, and provide a unique probe of magnetic fields associated with the ``central engine\" powering jets and outflows;  $\\bullet$ X-ray activity, down to very low levels, has an important impact (via irradiation processes) on physical conditions in protoplanetary disks, in particular ionization and coupling of material with magnetic fields, and by inference on the early solar system where products of internal spallation reactions may be recorded today in anomalous meteoritic abundances;  $\\bullet$ The increase in sensitivity and/or angular resolution has enlarged the number of observable mechanisms for YSO X-ray emission, to include shocks (both from accretion in T Tauri stars, and from jets in Herbig-Haro objects), and shock-related pc-scale diffuse emission processes from stellar winds of massive stars in HII regions.  At the same time, the limitations of \\Ch\\ and \\XMM\\ already begin to be visible:  for instance, high-resolution spectroscopy at 6.4 keV ($\\Delta E << 100$ eV), combined with a higher throughput to allow time resolution, is necessary to use ``reverberation mapping'' to probe the circumstellar disks of YSOs.  So, even as \\Ch\\ and \\XMM\\ results on star-forming regions keep coming, it is already time to follow the next-generation X-ray projects like {\\sl Astro-E II}, {\\sl XEUS} and {\\sl Constellation-X} !"
    },
    "0206/astro-ph0206496_arXiv.txt": {
        "abstract": "Strong gravitational lensing has traditionally been one of the few phenomena said to oppose a large cosmological constant; many analyses of lens statistics have given upper limits on $\\Omega_\\Lambda$ that are marginally inconsistent with the concordance cosmology. Those conclusions were based on models where the predicted number counts of galaxies at moderate redshifts ($z \\sim 0.5$--1) increased significantly with $\\Omega_\\Lambda$.  I argue that the models should now be calibrated by counts of distant galaxies. When this is done lens statistics lose most of their sensitivity to the cosmological constant. ",
        "introduction": "Popular opinion seems to have settled on a ``concordance'' cosmology dominated by dark energy. The cosmic microwave background indicates a flat geometry (e.g., de Bernardis et al.\\ 2000; Hanany et al.\\ 2000; Pryke et al.\\ 2002), cluster mass-to-light ratios indicate a low matter density (e.g., Carlberg, Yee \\& Ellingson 1997; Bahcall et al.\\ 2000), and type Ia supernovae indicate cosmic acceleration (e.g., Riess et al.\\ 1998; Perlmutter et al.\\ 1999). The popular cosmology has matter content $\\Omega_M \\approx 0.30$--0.35 and dark energy $\\Omega_X \\approx 0.65$--0.70 such that $\\Omega_M + \\Omega_X = 1$. For simplicity I assume that the dark energy corresponds to a cosmological constant $\\Lambda$, although my analysis could easily be extended to quintessence models (e.g., Waga \\& Miceli 1999).  One phenomenon that has traditionally stood out from the concordance model is strong gravitational lensing. The statistics of strong lenses are sensitive to the cosmological parameters via the cosmological volume element (e.g., Turner 1990; Fukugita, Futamase \\& Kasai 1990), and analyses of the data have yielded upper limits on the dark energy at the level of $\\Omega_\\Lambda < 0.66$ at 95\\% confidence (e.g., Kochanek 1996; Falco, Kochanek \\& Mu\\~noz 1998). There are small systematic uncertainties in the upper limit due to assumptions about the lens sample and the amount of dust extinction in lens galaxies (e.g., Falco et al.\\ 1998; Helbig et al.\\ 1999; Waga \\& Miceli 1999; Cooray, Quashnock \\& Miller 1999). Larger systematic effects arise from uncertainties in the local luminosity function of galaxies. Chiba \\& Yoshii (1999) argue that by adopting luminosity functions from different surveys they can relax the upper limit on $\\Omega_\\Lambda$ and even find models that favor values in the range $\\Omega_\\Lambda \\sim 0.5$--0.8. Kochanek et al.\\ (1998) respond by acknowledging the systematic uncertainties but defending their choice of luminosity function as the one that is most consistent with the observed luminosities of lens galaxies. This controversy will soon be resolved by new measurements of the luminosity function from the SDSS and 2dF surveys that appear to eliminate most of the traditional problems (Blanton et al.\\ 2001; Cross et al.\\ 2001).  While debating the details, the previous studies agreed on the idea that raising $\\Omega_\\Lambda$ dramatically increases the expected number of lenses; they agreed on the trend and only contested the zero point. In this {\\it Letter\\/} I question the trend itself, based on number counts of distant galaxies. I argue that the lens statistics models used to obtain upper limits on $\\Omega_\\Lambda$ are inconsistent with galaxy counts at $z \\sim 0.5$ --- not with any particular value of the counts, but with the general idea that they can be measured and used as constraints on the models. I modify the models to be calibrated by galaxy counts and consider how the new models depend on $\\Omega_\\Lambda$. For simplicity I consider only flat cosmologies. ",
        "conclusions": "The optical depth for lensing is basically proportional to the number of galaxies on the sky at redshifts $z_l \\sim 0.3$--1. Traditionally that number was not well known, so models for lens statistics adopted a number density of galaxies and multiplied by the cosmological volume to get the number. With the assumed number density held fixed, the expected number of lenses was proportional to the cosmological volume and hence very sensitive to $\\Omega_\\Lambda$.  Number counts of distant galaxies can now be measured directly, and they are inconsistent with the idea that the number density is independent of $\\Omega_\\Lambda$. This general point holds whether there is much or little redshift evolution in the galaxy population. Constraining models for lens statistics to agree on distant galaxy counts makes them far less sensitive to $\\Omega_\\Lambda$. Whereas the old models saw a factor of 3 change in the number of lenses between $\\Omega_\\Lambda=0$ and $\\Omega_\\Lambda=0.7$, in the new models the change is $\\lesssim$10--40\\%. Using lens statistics to constrain $\\Omega_\\Lambda$ may still be possible, but will be difficult.  This paper has focused on models where the deflector population is derived from observed galaxy populations, which I refer to as phenomenology models. In an alternate class that I call theory models, the deflector population is described with a mass function from structure formation theory; the resulting models are sensitive to cosmology not only through the volume element but also through $\\Omega_M$ and the growth of structure. In theory models the predicted number of lenses {\\it decreases\\/} as $\\Omega_\\Lambda$ increases (for flat cosmologies; Porciani \\& Madau 2000; Li \\& Ostriker 2002), which strongly disagrees with old phenomenology models but is less different from my new models. In quintessence models, phenomenology and theory models do agree that making the equation of state more negative increases the predicted number of lenses (Waga \\& Miceli 1999; Sarbu, Rusin \\& Ma 2001). The differences between phenomenology and theory models clearly need further study. Nevertheless, I would argue that the focus of lens statistics should move away from constraining $\\Omega_\\Lambda$ and toward learning about the population of dark matter halos out to $z \\sim 1$."
    },
    "0206/astro-ph0206175_arXiv.txt": {
        "abstract": "We have carried out a kinematical, high angular resolution ($\\sim$ 0\\farcs 1) study of the optical blueshifted flow from DG~Tau within 0\\farcs 5 from the source (i.e.\\ 110 AU when de-projected along this flow). We analysed optical emission line profiles extracted from a set of seven long-slit spectra taken with the  {\\em Space Telescope Imaging Spectrograph} (STIS) on board the Hubble Space Telescope (HST), obtained by maintaining the slit parallel to the outflow axis while at the same time moving it transversely in steps of 0\\farcs 07. For the spatially resolved flow of moderate velocity (peaking at -70\\,km\\,s$^{-1}$), we have found systematic differences in the radial velocities of lines from opposing slit positions i.e.\\ on alternate sides of the jet axis. The results, obtained using two independent techniques, are corrected for the spurious wavelength shift due to the uneven illumination of the STIS slit. Other instrumental effects are shown to be either absent or unimportant. The derived relative Doppler shifts range from 5 to 20\\,km\\,s$^{-1}$. Assuming the flow is axially symmetric, the velocity shifts are consistent with the southeastern side of the flow moving towards the observer faster than the corresponding northwestern side. If this finding is interpreted as rotation, the flow is then rotating clockwise looking from the jet towards the source and the derived toroidal velocities are in the range  6 to 15\\,km\\,s$^{-1}$, depending on position. Combining these values with recent estimates of the mass loss rate, one would obtain an angular momentum flux, for the low to moderate velocity regime of the flow, of $\\dot{J}_{w,lm}\\sim 3.8\\,10^{-5}$\\,M$_\\odot$\\,yr$^{-1}$ AU\\,km\\,s$^{-1}$. Our findings may constitute the first detection of rotation in the initial channel of a jet flow. The derived values appear to be consistent with the predictions of popular magneto-centrifugal jet-launching models, although we cannot exclude the possibility that the observed velocity differences are due to some transverse outflow asymmetry other than rotation. ",
        "introduction": "\\label{intro}  Herbig-Haro (HH) jets, optically emitting collimated mass outflows associated with young stellar objects (see, e.g., \\citet{ray98}, \\citet{eis00}), are widely recognised as an essential ingredient of the star formation process. In particular, they are believed to contribute to the removal of excess angular momentum from accreted matter and to disperse infalling circumstellar envelopes. Despite their possible key role in star formation, the origin of the jets themselves remains elusive, although it is believed that their generation involves the simultaneous action of magnetic and centrifugal forces in a rotating star/disk system \\citep{konigl00,shu00,shib99}. Canonical models, however, have not been tested observationally, since the process is believed to occur on very small scales (less than a few AU), although according to some, the acceleration and collimation region may extend to $\\sim$\\,100\\,AU from the star.  With these ideas in mind, we have observed  DG~Tau with STIS on-board the HST. Multiple overlapping slit positions parallel to the outflow from this star were chosen so as to  build up a 3-D spatial intensity-velocity ``cube'' in various forbidden emission lines (FELs) and H$\\alpha$. In a previous paper \\citep{bmresc00} we have presented high spatial resolution ($\\sim$ 0\\farcs 1) velocity channel maps of the first 2\\arcsec\\ of this flow in different emission lines and in several distinct radial velocity intervals. In a subsequent paper \\citep{brmes02} we will present mono- and bi-dimensional maps of the excitation and dynamical properties of the same region of the flow, also in various velocity intervals. Briefly the main conclusions of these studies are as follows. The outflow appears to have an onion-like kinematic structure, with the faster and more collimated gas continuously bracketed by wider and slower material. In addition, the flow becomes gradually denser, and of higher excitation, with proximity to the central axis. Combining these results, we have been able to calculate a flow mass loss rate, $\\dot{M}_w \\sim 2.4~10^{-7}$\\,M$_\\odot$\\,yr$^{-1}$, which is about one tenth of the estimated mass accretion rate through the disk \\citep{heg95,brmes02}. These results are in line with what is predicted by popular magneto-hydrodynamic (MHD) models.  Further evidence, however, that would help to support such models is observation of rotation in the initial section of the flow. In the magneto-centrifugal scenario, the flow maintains a record of rotation at its base, at least initially, during propagation. If the flow has a favourable inclination angle with respect to the line of sight, a trace of its rotation should be seen in sets of high angular resolution spectra taken close to the source with the slit parallel to the outflow axis. Our DG~Tau dataset is ideally suited for such an investigation. We note that hints of rotation were found in the HH~212 flow at large distances (2\\,10$^3$ to 10$^4$\\,AU) from the source by \\citet{dbsch00}. Although these results are very encouraging, we nevertheless believe that the launching mechanism is better constrained by the kinematical properties of the {\\em initial portion} of the jet channel. Here, the flow may not have suffered the effects of strong interactions with its environment.  The results of our analysis are described in Section~\\ref{rot}, after a brief summary of our observations and data reduction (Section~\\ref{obs}). We then discuss our findings in Section~\\ref{disc} in the light of MHD jet models. ",
        "conclusions": "\\label{conc}  Many features of the collimated HH jets associated with star formation are still unexplained, especially the origin of the jets themselves. Models for the launching of flows still lack observational constraints, due to the small size of the acceleration zone and the fact that the central sources are often heavily embedded. In order to shed some light on this area, we have taken and analysed a set of seven high angular resolution  ($\\leq 0\\farcs1$) HST/STIS spectra of the outflow from the T Tauri star DG~Tau. Here we report possible evidence for rotation in this dataset: rotation, in fact, is a fundamental ingredient for the modelling of the acceleration of outflows and of the interplay between accretion and ejection of matter in the framework of the formation of a new star.  From a detailed analysis of the line profiles in each spectrum, and in four distinct regions of the initial part of the jet (within about 100\\,AU from the star), we have found systematic differences in the radial velocity of the lines for each pair of slits displaced symmetrically with respect to the jet axis. The values, obtained with multiple Gaussian fitting and/or cross-correlation routines, have been  corrected for the wavelength offset produced by the uneven illumination of the STIS slit. Other possible instrumental effects that may contaminate the data are either absent or unimportant. According to our results, and under the assumption that the flow is axially symmetric, the southeastern side of the blueshifted jet appears to move toward the observer faster than the corresponding northwestern side, and the average value found for the difference is about 10\\,km\\,s$^{-1}$. A detailed map of the shifts is available in Fig.\\,\\ref{spinmap}, derived from the values in Table\\,\\ref{tbl-1}. If we interpret these findings in terms of rotation of the flow, they would imply that the jet is rotating clockwise looking from the flow tip towards the source. Taking into account the inclination of this system with respect to the line of sight, one would derive apparent toroidal velocities of around 6 to 15\\,km\\,s$^{-1}$ at a few tens of AU from the axis, and between 20 to 90\\,AU above the disk plane. These velocities would be in the range predicted by MHD theories. We also estimate  the angular momentum flux in the low to medium velocity regime to be about 3.8\\,10$^{-5}$\\,M$_\\odot$\\,yr$^{-1}$\\,AU\\,km\\,s$^{-1}$. This could amount to 60\\% of the angular momentum that the disk has to loose per unit time at the footpoint in order to accrete at the observed rate. Our findings may constitute the first detection of rotation in the initial channel of a jet flow, although at the present stage we cannot exclude that non-uniformities in the jet and/or in its environment cause  the observed asymmetries. More observations of this and other outflows are required. If the future studies will confirm the presence of rotational motions at the base of the flows, it would represent an important validation of magneto-centrifugal models for the launching of both galactic and extra-galactic jets."
    },
    "0206/gr-qc0206022_arXiv.txt": {
        "abstract": "{A relativistic sub-picosecond model of gravitational time delay in radio astronomical observations is worked out and a new experimental test of general relativity is discussed in which the effect of retardation of gravity associated with its finite speed can be observed.  As a consequence, the speed of gravity can be measured by differential VLBI observations. Retardation in propagation of gravity is a central part of the Einstein theory of general relativity which has not been tested directly so far. The idea of the proposed gravitational experiment is based on the fact that gravity in general relativity propagates with finite speed so that the deflection of light caused by the body must be sensitive to the ratio of the body's velocity to the speed of gravity. The interferometric experiment can be performed, for example, during the very close angular passage of a quasar by Jupiter. Due to the finite speed of gravity and orbital motion of Jupiter, the variation in its gravitational field reaches observer on Earth not instantaneously but at the retarded instant of time and should appear as a velocity-dependent excess time delay in addition to the well-known Shapiro delay, caused by the static part of the Jupiter's gravitational field. Such Jupiter-QSO encounter events take place once in a decade. The next such event will occur on September 8, 2002 when Jupiter will pass by quasar J0842+1835 at the angular distance $3.7'$. If radio interferometric measurement of the quasar coordinates in the sky are done with the precision of a few picoseconds ($\\sim$ 5 $\\mu$as) the effect of retardation of gravity and its speed of propagation may be measured with an accuracy about 10\\%.} \\date{} \\titlerunning{Relativistic model for experimental measurement of the speed of gravity } ",
        "introduction": " ",
        "conclusions": "We believe that the differential VLBI experiment in September 2002 can measure the retardation effect in propagation of gravity and determine the speed $c_g$ of its propagation with 10\\% to 20\\% accuracy.  If the experiment is successful it will provide a new independent test of general relativity in the solar system."
    },
    "0206/astro-ph0206065_arXiv.txt": {
        "abstract": "The massive X-ray binary Cen~X-3 was observed over approximately one quarter of the system's 2.08 day orbit, beginning before eclipse and ending slightly after eclipse center with the {\\it Chandra X-ray Observatory} using its High-Energy Transmission Grating Spectrometer. The spectra show K shell emission lines from hydrogen- and helium-like ions of magnesium, silicon, sulfur, and iron as well as a K$\\alpha$ fluorescence emission feature from near-neutral iron.  The helium-like $n=2\\to1$ triplet of silicon is fully resolved and the analogous triplet of iron is partially resolved.  We measure fluxes, shifts, and widths of the observed emission lines.  The helium-like triplet component flux ratios outside of eclipse are consistent with emission from recombination and subsequent cascades (recombination radiation) from a photoionized plasma with temperature $\\sim$100\\,eV.  In eclipse, however, the $w$ (resonance) lines of silicon and iron are stronger than that expected for recombination radiation, and are consistent with emission from a collisionally ionized plasma with a temperature of $\\sim$1\\,keV.  The triplet line flux ratios at both phases can be explained more naturally, however, as emission from a photoionized plasma if the effects of resonant line scattering, a process which has generally been neglected in X-ray spectroscopy, are included in addition to recombination radiation.  We show that resonant line scattering in photoionized plasmas may increase the emissivity of $n=2\\to1$ line emission in hydrogen and helium-like ions by a factor as large as four relative to that of pure recombination and so previous studies, in which resonant scattering has been neglected, may contain significant errors in the derived plasma parameters.  The emissivity due to resonance scattering depends sensitively on the line optical depth and, in the case of winds in X-ray binaries, this allows constraints on the wind velocity even when Doppler shifts cannot be resolved. ",
        "introduction": "\\label{sec:intro}  It has long been known that in eclipsing high mass X-ray binaries (HMXBs, X-ray binaries where the companion star is a massive star of type O or B) a residual X-ray flux can be observed during eclipse of the compact X-ray source \\citep{sch72}.  \\citet{bec78} observed that in Vela~X-1, the spectral shape of the hard X-ray flux observed during eclipse was approximately the same as that observed out of eclipse, and inferred from this that the residual eclipse flux was due to electron scattering of the primary X-rays in circumstellar material. The residual eclipse fluxes of X-ray binaries are typically a few percent of the out-of-eclipse fluxes, implying that the scattering optical depth in the continuum is approximately a few percent.  Using this electron scattering optical depth and, for the length scale, the orbital separation $a$ (of order $10^{12}$\\,cm) circumstellar densities of order $10^{10}$\\,cm$^{-3}$ are inferred, assuming $\\tau=n_e\\sigma_{\\rm T}a$, where $\\sigma_{\\rm T}$ is the Thomson cross-section.  Isolated O and B type stars have winds with mass-loss rates of order $10^{-7}$--$10^{-6}\\,M_\\sun$\\,yr$^{-1}$ and velocities of order 1000\\,\\kms.  The winds are driven by ultraviolet photons which impart their outward momentum to the matter primarily through scattering in line transitions \\citep{luc70,cas75}.  For a spherically symmetric steady-state wind, the hydrogen atom density ($n$) is related to the mass-loss rate ($\\mdot$) and velocity ($v$) by \\begin{equation} n=\\frac{\\mdot}{4\\pi\\mump r^2v} \\label{eqn:n} \\end{equation} where $r$ is the distance from the stellar center, $m_p$ is the proton mass, and we take $\\mu$, the mass in amu per hydrogen atom, to be 1.4. Using, again, the length scale $10^{12}$\\,cm, the average density derived above from electron scattering is consistent with the densities in the winds of massive stars.  This similarity suggests that the dynamics of hot star winds in X-ray binary systems might be quite similar to the those in isolated stars. However, the X-ray luminosities ($L$) of these objects are of order $10^{37}$\\,\\ergs, implying that typical values of the ionization parameter \\citep[$\\xi=L/ns^2$, where $s$ is the distance from the X-ray source,][]{tar69} in the circumstellar material are of order $10^3$\\,\\ergs{}cm, meaning that the wind should be highly ionized.  At such high ionization, the UV line opacity of the wind is greatly reduced and, therefore, so is the force exerted on the wind by the star's radiation \\citep{ste90}.  Several calculations (which, however, have several significant approximations) have shown that UV photons cannot drive a wind on the X-ray illuminated side of an X-ray luminous HMXB \\citep[e.g.,][]{ste91}.  It may be that an alternative mechanism, such as evaporation from the X-ray heated surface of the companion \\citep[e.g.,][]{bas77} drives the wind in luminous HMXBs.  The high ionization conditions in the winds of HMXBs can be inferred directly by the observation of K shell emission lines from hydrogen- and helium-like ions.  \\citet{nag92}, for example, noticed in a \\ginga\\ observation of Cen~X-3 that the iron line was not consistent with only a 6.4\\,keV line from neutral iron but could be fit if a 6.7\\,keV component from helium-like iron was included. The launch of the \\asca\\ observatory, with its Solid-State Imaging Spectrometers (SIS), represented a great improvement over previous observatories in the sensitivity to X-ray emission lines.  With the SIS instrument, winds of HMXBs were seen to produce K-shell emission lines from hydrogen-like, helium-like, and near-neutral ions of elements from neon through iron \\citep{nag94,ebi96}.  For $\\xi\\sim10^3$\\,\\ergs{}cm, recombination and subsequent electronic cascades (recombination radiation) are an efficient source of line emission.  In several studies, X-ray emission line spectra observed with \\asca\\ were used to characterize the winds in HMXBs using line emissivities due to recombination.  Two of these systems, Vela~X-1 and Cen~X-3, were found to have very different wind characteristics that may serve to illuminate important physical processes which occur in X-ray binary winds.  The characteristics of the wind in Vela~X-1 were explored by \\citet{sak99} by fitting the emission-line spectrum observed during eclipse to the recombination radiation calculated from model winds. It was found that the emission lines from hydrogen and helium-like lines could be fit by a wind model with a density (or equivalently, for a given wind velocity, a mass-loss rate) approximately an order of magnitude smaller than that inferred in other other studies using different methods \\citep[e.g.,][using the electron scattered continuum as discussed above, and the 6.4\\,keV iron fluorescence line]{lew92}. The emission line spectrum also exhibited fluorescence lines from near-neutral ions.  All of the observational data could be explained if the wind had a population of dense clumps filling only a small part of the wind volume but containing most ($\\sim$90\\%) of the mass of the wind.  It is conceivable that a wind could be driven by radiation pressure on the low-ionization material in the clumps.  By contrast, Cen~X-3 has an X-ray luminosity of $10^{38}$\\,\\ergs, more than an order of magnitude greater than that of Vela~X-1.  Its spectrum, as seen by \\asca, does not contain any fluorescence lines produced in the wind\\footnote{A fluorescent iron line feature observed from Cen~X-3 is produced near the neutron star \\citep{nag92,day93b,ebi96}}, indicating the absence of a low-ionization component.  \\citet{woj01} showed that its emission line spectrum can be explained by recombination from a smooth, highly ionized wind.  It was speculated that the large X-ray flux in the wind of Cen~X-3 might evaporate clumps or prohibit them from forming.  Radiation driving of such a smooth, highly ionized wind does not appear plausible, and it would appear that an alternate mechanism, such as X-ray heating of the companion is necessary \\citep[e.g.,][]{day93a}.  The winds of HMXBs have also been studied using X-ray emission lines and the assumption that the emission lines are due to recombination radiation by \\citet{ebi96} and \\citet{bor01}.  For a given X-ray luminosity and a given system geometry, the luminosity of recombination lines from the wind of a HMXB depends on the wind parameters only through its density distribution, as the temperature and ionization state are determined by the ionization parameter.  As can be seen from Equation~\\ref{eqn:n}, the density, and therefore the recombination line emissivity, depends on the mass-loss rate and the velocity only in the combination $\\mdot/v$. \\citet{woj01} has described explicitly how, for a velocity profile which scales with the parameter $\\vinf$ (the terminal velocity), the luminosity of the recombination lines depends on $\\mdot$ and $\\vinf$ only in the combination $\\mdot/\\vinf$, and derived for Cen~X-3 a value of $\\mdot/\\vinf\\approx10^{-6}\\,M_\\sun$yr$^{-1}(1000\\,{\\rm km\\,s^{-1}})^{-1}$.  The degeneracy in these parameters could be removed if velocities could be measured independently, such as through Doppler line shifts.  The resolution of the SIS detectors was approximately 2\\% in the iron K region (6.4--7.0\\,keV), and worse at the lower energies of the other emission lines, and so was insufficient to detect Doppler shifts of less than $\\sim$2000\\,\\kms. As this is of the same order as the typical terminal velocity of UV driven winds, this has not allowed for strong constraints on the wind driving mechanism.  The gratings on \\chandra\\ however, make possible detection of Doppler shifts as small as a few hundred \\kms\\ and therefore allow, in principle, tighter constraints on the wind driving mechanism.  The HMXB Cen~X-3 is the most luminous persistent HMXB in the galaxy and, presumably, is an extreme example of disruption of the wind by X-rays.  It is, therefore, a strong candidate for alternative wind driving mechanisms.  We observed Cen~X-3 over an eclipse with the \\chandra\\ HETGS in order to constrain the wind velocity by resolving emission line Doppler shifts and thus constrain the mass-loss rate and possible wind driving mechanisms.  The spectra we obtain from our observations exhibit emission lines from hydrogen- and helium-like ions from magnesium to iron, as well as a a fluorescent line from near-neutral iron.  We resolve the $n=2\\to1$ triplet of helium-like silicon, partially resolve the helium-like triplet of iron, and derive constraints on the respective flux ratios of the triplet components.  Outside of eclipse, the constraints on the ratios are consistent with recombination radiation.  However, in eclipse, in contrast to our expectations, the constraints are {\\em not\\/} consistent with those expected for recombination.  The $w$ (resonance) lines are more intense, relative to the other components of the triplets, than expected for recombination and are, in fact, consistent with emission from a collisionally ionized plasma.  Similar enhanced $w$ lines have also been seen from Vela~X-1 \\citep{sch02}.  This, at first, appears to be a significant challenge to our paradigm of HMXB winds.  Hydrogen- and helium-like line emission in collisionally ionized plasmas occurs at much higher temperatures ($\\sim$1\\,keV) than in photoionized plasmas ($\\sim$100\\,eV) and a mechanical source of heating, such as shocks in the wind, would be required to maintain such a high temperature.  However, while contributions to the line spectrum from collisionally-ionized gas cannot be ruled out a priori, it is more natural to attribute the enhanced fluxes of lines with large oscillator strengths to direct photoexcitation by radiation from the neutron star.  Indeed, the dangers of neglecting photoexcitation have long ago been pointed out for the case of optical spectra of planetary nebulae \\citep{sea68}, in which a hot star photoionizes a cloud of surrounding gas.  More recently, clear evidence for photoexcitation-driven line emission has also been obtained from Seyfert 2 spectra, where similarly enhanced He-like $w$ lines have been observed with \\chandra\\ HETGS  \\citep{sak00b}.  The presence of absorption lines in Seyfert 1 galaxies (where a direct line of sight to the compact radiation source exists) gives further evidence for photoexcitation in these objects \\citep{kaa00,kas01}. X-ray spectra of Seyfert galaxies are similar to those of HMXBs, implying that plasma in the two types of systems has similar ionization conditions ($\\log\\xi\\sim$2--4).  In both types of systems, plasma column densities are of order $10^{21}$\\,cm$^{-2}$. Furthermore, in both types of systems, the X-ray emitting gas is thought to have bulk motions with velocities of order 100 to 1000\\,\\kms.  Therefore, we should not be surprised that photoexcitation is an important process in producing the observed spectra in HMXBs.  As resonant X-ray line scattering has received little attention in the literature, in this paper we focus on the effects of resonant scattering in HMXBs, with the expectation that this work will complement parallel studies in the Seyfert galaxy domain.  We show that resonant scattering of radiation from the compact source in an HMXB results in increased line fluxes for lines with large oscillator strengths (such as the $w$ lines of helium-like ions) in observations during X-ray eclipse.  We further show that this line flux enhancement due to resonant scattering saturates as line optical depths become comparable to unity and that the relative line fluxes that we observe require non-zero line optical depths.  The line optical depths along a path depend on the velocity distribution along that path.  \\citet{kin02} have calculated the saturation of the line flux enhancement due to resonant scattering in plasmas without bulk motions but with a Gaussian distribution of ion velocities and shown that the \\xmm\\ RGS emission line spectrum of the Seyfert 2 galaxy NGC~1068 can be fit for appropriate values of the ion column density.  In a gas where bulk motions are larger than the thermal or other small-scale velocities, the optical depth, and therefore, the depletion of the resonantly scattered line luminosity, depends on the bulk velocity distribution. In the context of an HMXB wind, therefore, the resonantly scattered line luminosity has a dependence on $\\vinf$ other than in the combination $\\mdot/\\vinf$.  We show that in spectra where the helium-like $n=2\\to1$ triplets can be resolved, the portion of line luminosity due to resonant scattering may be derived, and, for a given wind model, explicit constraints on the terminal wind velocity may be derived.  These constraints may be derived only from observed line fluxes, and do not require any information regarding Doppler line shifts or broadening.  In \\S\\ref{sec:obs}, we describe our observations and reduction of the data.  In \\S\\ref{sec:spec_anal_lines}, we fit the line fluxes, widths, and shifts and derive constraints on the ratios of the fluxes of the components of the helium-like triplets.  In \\S\\ref{sec:res_scat} we calculate line emissivities due to resonant scattering, show that the observed line ratios can be explained by resonant scattering, and describe how constraints on the wind velocity can be derived. In \\S\\ref{sec:discussion} we summarize our conclusions and discuss the errors inherent in the previous studies which have not included the effects of resonant scattering. ",
        "conclusions": "\\label{sec:discussion}  We have observed Cen~X-3 from before an eclipse until mid-eclipse. During this observation, the continuum flux underwent a large increase, then decreased due to the eclipse ingress.  We have fit the shifts, widths, and fluxes of observable lines.  Our best fit velocity shifts and (Gaussian $\\sigma$) velocity widths are generally less than 500\\,\\kms.  These velocities are significantly smaller than terminal velocities of isolated O star winds \\citep[1--2$\\ee3$\\,\\kms, e.g.,][]{lam99}.  However, in isolated O stars, X-ray line velocity widths (HWHM or Gaussian $\\sigma$) have been observed from $\\sim$400\\,\\kms \\citep{sch00} to $\\sim1000$\\,\\kms\\ \\citep{wal01,cas01,kah01a}.  Our results for the Doppler velocities are therefore consistent with those in isolated O stars.  We have measured the ratios of the fluxes of the components of the helium-like triplets of silicon and iron.  We find that the flux ratios are consistent with recombination in a photoionized plasma, with the exception that during the eclipse phase, the $w$ line is stronger than expected from recombination.  Resonant scattering provides a natural mechanism for increasing the fluxes of the $w$ lines in eclipse but not out of it.  We have calculated the enhancement of the $w$ line fluxes due to resonant scattering and we find that the observed relative $w$ line fluxes are smaller than expected for resonant scattering in the optically thin limit and non-zero line optical depths are required.  We are therefore able to explain the observed emission line triplets of helium-like ions --- which, during eclipse, appear very much like what would be observed from a hot, collisionally ionized plasma --- without rejecting the hypothesis that the wind is composed entirely of a warm photoionionized plasma.  As mentioned in \\S\\ref{sec:intro}, in previous spectral studies of HMXB winds with moderate-resolution spectral data, X-ray emission line spectra have been interpreted with the assumption that line emission is due purely to recombination radiation.  The primary physical quantity inferred from a line luminosity is, as the emissivity due to recombination depends on the square of the density, an emission measure: $\\int n_e^2dV$.  If the apparent line emissivity is underestimated, such as by neglect of resonant scattering, then the inferred emission measure is overestimated by that same amount.  The parameter $\\mdot/\\vinf$ is proportional to the wind density and therefore proportional to the square root of the wind emission measure.  We have shown that outside of eclipse, resonant scattering along the line of sight to the neutron star offsets the line flux enhancement due to resonant scattering in the bulk of the gas.  In our case, however, these two effects appear to have may nearly cancelled, and so for interpreting spectra obtained outside of eclipse, it may be a good approximation to ignore resonant scattering.  For spectra obtained during eclipse, however, resonant scattering, in the optically thin limit, increases the $n=2\\to1$ line fluxes of hydrogen- and helium-like ions by factors of, respectively 3.9 and 2.3. Therefore, emission measures inferred assuming line emissivities from recombination only, from spectra obtained during eclipse, could be too large by factors as large as four.  However, in fact, the error due to neglect of resonant scattering is probably not so large as the winds may have significant optical depths in the lines, which, as discussed earlier, decreases the line luminosities due to resonant scattering relative to the line luminosities due to recombination.  Indeed, as we have shown in \\S\\ref{sec:res_scat}, for the eclipse spectrum described here, the total $n=2\\to1$ flux due to resonant scattering for \\ion{Si}{13} and \\ion{Fe}{25} is only approximately 0.6 and 1.0, respectively that due to recombination.  Therefore, previous estimates of the parameter $\\mdot/\\vinf$ may be too large only by factors of about two or less.  While this does not qualitatively the affect conclusions reached in any previous studies, the effects of resonant scattering should be included, even in the analysis of spectra of medium-resolution spectra, if only so the basic wind parameters may be derived accurately.  While the neglect of resonant scattering in previous work may not have resulted in significant errors in derived wind parameters, interpretation of high-resolution X-ray spectra with proper consideration of resonant scattering allows the determination of wind parameters that cannot otherwise be determined.  As \\citet{woj01} have shown, for a typical model wind, the wind density, and therefore the luminosity of recombination line emission depends on the parameters $\\mdot$ (the mass-loss rate) and $\\vinf$ (the terminal wind velocity) only in the combination $\\mdot/\\vinf$ and that it is therefore impossible to constrain either of these parameters individually using only the observed luminosities of lines resulting from recombination. Resonant scattering, however, has a different dependence on these parameters.  Emission from resonant scattering results as photons from the compact object are redirected by scattering ions toward the observer.  We imagine a model wind for which we decrease the velocity parameter $\\vinf$ while keeping the value of $\\mdot/\\vinf$, and therefore the density, constant.  As $\\vinf$ is decreased, the Doppler shifts decrease and the same number of ions scatter photons from a decreasing frequency range.  This results in a decrease in the number of photons which are subject to scattering but an increase in the scattering optical depth and an increase in the fraction of the accessible photons which can be scattered.  However, the fraction of the accessible photons which may be scattered is, of course, limited to be no more than unity and as $\\vinf$ is decreased, the total resonantly scattered line luminosity decreases.  Because $\\mdot/\\vinf$ is kept constant, however, the line luminosity due to recombination remains constant.  The emission line luminosities due to recombination and resonant scattering, therefore, depend on $\\mdot$ and $\\vinf$ in a non-degenerate way.  If the line luminosities due to recombination and resonant scattering can be determined, independent constraints on $\\mdot$ and $\\vinf$ can be derived.  We have demonstrated that, with high resolution spectroscopy, it is possible to resolve helium-like triplets and thereby discriminite between the line emission due to recombination and resonant scattering.  Therefore, using high resolution X-ray spectroscopic data, it is possible to derive independent constraints on the parameters $\\mdot$ and $\\vinf$. However, deriving these constraints requires detailed modeling of the wind, which is beyond the scope of this work, so we do not derive such constraints here.  However, this effect is most sensitive in the regime where the line optical depth is of order unity.  As we have shown in \\S~\\ref{sec:obs_const}, the line optical depths of the $w$ lines of helium-like silicon and iron in Cen~X-3 are, in fact, of order unity and so the prospects for using this effect to derive independent constraints on the wind parameters $\\mdot$ and $\\vinf$ are good.  Furthermore, these constraints do not depend at all on resolving Doppler line shifts or broadenings and are therefore independent of any constraints which may be obtained from observed line profiles.  From our analysis of \\chandra\\ spectra of Cen~X-3, we may conclude that in the analysis of high-resolution spectroscopic data, such as is obtained using the gratings on \\chandra\\ and \\xmm, of the winds of X-ray binaries and X-ray photoionized plasmas in general, it is critical that the effects of resonant scattering be included, if observed spectra are to be reproduced, or if inferences are to be made regarding the conditions of the emitting plasma (e.g., for determining whether plasma is photoionized, collisionally ionized, or ionized by a ``hybrid'' of both processes).  Furthermore, in the analysis of high-resolution spectroscopic data from HMXBs, including the effects of resonant scattering not only allows wind parameters to be derived accurately, but allows new constraints on wind parameters."
    },
    "0206/astro-ph0206253_arXiv.txt": {
        "abstract": "We present the results of a spectroscopic monitoring program (from 1998 to 2002) of the H$\\alpha$ emission strength in HDE 226868, the optical counterpart of the black hole binary, Cyg X-1.   The feature provides an important probe of the mass loss rate in the base of the stellar wind of the supergiant star.   We derive an updated ephemeris for the orbit based upon radial velocities measured from \\ion{He}{1} $\\lambda 6678$. We list net equivalent widths for the entire H$\\alpha$ emission/absorption complex, and we find that there are large variations in emission strength over both long (years) and short (hours to days) time spans. There are coherent orbital phase related variations in the profiles when the spectra are grouped by H$\\alpha$ equivalent width.  The profiles consist of (1) a P~Cygni component associated with the wind of the supergiant, (2) emission components that attain high velocity at the conjunctions and that probably form in enhanced outflows both towards and away from the black hole, and (3) an emission component that moves in anti-phase with the supergiant's motion.   We argue that the third component forms in accreted gas near the black hole, and the radial velocity curve of the emission is consistent with a mass ratio of $M_{\\rm X}/M_{\\rm opt} \\approx 0.36\\pm0.05$.   We find that there is a general anti-correlation between the H$\\alpha$ emission strength and X-ray flux (from the {\\it Rossi X-ray Timing Explorer} All Sky Monitor instrument) in the sense that when the H$\\alpha$ emission is strong ($W_\\lambda < -0.5$ \\AA ) the X-ray flux is weaker and the spectrum harder.   On the other hand, there is no correlation between H$\\alpha$ emission strength and X-ray flux when H$\\alpha$ is weak.   We argue that this relationship is not caused by wind X-ray absorption nor by the reduction in H$\\alpha$ emissivity by X-ray heating. Instead, we suggest that the H$\\alpha$ variations track changes in wind density and strength near the photosphere. The density of the wind determines the size of X-ray ionization zones surrounding the black hole, and these in turn control the acceleration of the wind in the direction of the black hole. During the low/hard X-ray state, the strong wind is fast and the accretion rate is relatively low, while in the high/soft state the weaker, highly ionized wind attains only a moderate velocity and the accretion rate increases.  We argue that the X-ray transitions from the normal low/hard to the rare high/soft state are triggered by episodes of decreased mass loss rate in the supergiant donor star. ",
        "introduction": "%  Cygnus~X-1 has been one of the most intensively studied X-ray sources in the sky since its discovery and identification with the O9.7 Iab supergiant star, HDE 226868 \\citep{bol72,web72}. This system provided the first evidence for the existence of stellar mass black holes when it was discovered to be a 5.6 day binary with a massive, unseen companion. \\citet{gie86a} used the spectroscopic orbit \\citep{gie82,las98,bro99a}, light curve \\citep{kem83,kar01}, photospheric line broadening, and a range in the assumed degree of Roche-filling of the supergiant star to obtain mass estimates of $M_{\\rm opt} = 23 - 43 ~M_\\odot$ and $M_{\\rm X} = 10 - 21 ~M_\\odot$.  \\citet{her95} derived physical parameters for the visible supergiant based upon a spectroscopic analysis of the line spectrum, and they adopted a system inclination of $i=35^\\circ$ (based upon published estimates) to arrive at mass estimates of 18 and $10 ~M_\\odot$ for the supergiant and black hole, respectively (the derived masses scale as $\\sin^{-3} i$ over the probable range of $i= 30^\\circ - 40^\\circ$; \\citet{gie86a,wen99}). The X-ray source in Cyg~X-1 is powered mainly by accretion from the strong stellar wind of the supergiant star \\citep{pet78,kap98a}. In fact, Cyg~X-1 probably represents a situation intermediate between pure, spherical wind accretion and accretion by Roche lobe overflow.  Observations of the optical emission lines \\citep{gie86b,nin87} indicate that the wind departs from spherical symmetry and that there exists an enhanced wind flow (or ``focused wind'') in the direction of the companion. The intense X-ray emission is believed to be produced close to the black hole in an accretion disk that emits soft X-ray photons and in a hot corona that inverse-Compton scatters low energy photons to higher energies \\citep{lia84,tan95}. Ultraviolet radiation from close to the black hole has been detected through High Speed Photometer observations with the {\\it Hubble Space Telescope} \\citep{dol01}. Radio jets were recently discovered in Cyg~X-1 \\citep{sti01,fen01} indicating a collimated outflow with a speed in excess of $0.6 c$, so that Cyg~X-1 joins the group of Galactic {\\it microquasars}, small scale versions of active galactic nuclei \\citep{mir99}.  The target is also a candidate $\\gamma$-ray transient source \\citep{gol02}; the $\\gamma$-rays are probably created through inverse-Compton scattering in the jets.  Cyg~X-1 is generally found in either a low/hard state (the more common case of low 2-10 keV flux and a hard energy spectrum) or a high/soft state (in which the soft X-ray flux increases dramatically and the spectrum softens; \\citet{zha97,zdz02}).  Every few years Cyg~X-1 makes a transition from the low/hard to the high/soft state, and it remains in this active state for weeks to months before returning to the low/hard state.  The last well documented high/soft state occurred in 1996 \\citep{bro99b}, and in 2001 September Cyg~X-1 once again entered a high/soft state in which it still remains at the time of writing (2002 September). This recent transition into the high/soft state was accompanied by a sudden decrease in radio flux \\citep{poo01} (the opposite of the radio brightening that accompanied the return to the low/hard state in 1971 and that led directly to the identification of the visible star associated with the X-ray source; \\citet{hje71}).  There are many theories about the causes of the transitions \\citep{cha95,pou97,mey00,wen01,you01,rob02} which generally relate to the physical conditions of the gas surrounding the black hole.   For example, \\citet{esi98} describe the transitions in terms of an advection-dominated accretion flow (ADAF) model in which the transitions are related to changes in the inner radius of the geometrically thin, optically thick, Keplerian disk. In the usual low/hard state, the inner disk radius is relatively large, but during the high/soft state the inner radius extends inwards close to the last stable orbit around the black hole.  The transition to the high/soft state is generally believed to be the result of a moderate increase in the mass accretion rate.  However, there is no clear observational evidence available to support the claim of enhanced mass transfer during the high/soft state. In fact, the evidence collected so far hints that the supergiant mass loss rate may actually decline during the high/soft state. \\citet{wen99} present a model for the X-ray light curve of Cyg~X-1 during the prolonged low/hard state based upon the accumulated data from the {\\it Rossi X-ray Timing Explorer Satellite} (RXTE) All Sky Monitor (ASM) instrument (see also \\citet{kar01}). They find that the decreased X-ray flux observed when the supergiant is in the foreground can be explained by the X-ray absorption caused by the wind outflow from the supergiant.  However, during the high/soft state the orbital variation in X-ray flux disappears, and Wen et al.\\ argue that the wind absorption declines because of increased photoionization of the wind by the stronger X-ray source and a decrease in the wind density (implying a factor of 2 decrease in the mass loss rate). \\citet{vol97} found that the H$\\alpha$ emission associated with the wind loss also declined during the 1996 high/soft state. Taken at face value, these observations suggest that the wind mass loss actually declines during a high/soft state, in contradiction to the theoretical expectations.  In this paper we report on multiple year observations of the H$\\alpha$ emission line in HDE 226868 (\\S2) which we find to show significant long term and short term variability (\\S3), presumably reflecting changes in the wind close to the supergiant.   The RXTE/ASM instrument has provided continuous X-ray flux measurements of the binary throughout this period, and we show that temporal variations in the X-ray flux are broadly anti-correlated with the H$\\alpha$ emission strength (\\S4).   We discuss the implications of this result, and we suggest that the state transitions may result from changes in the wind velocity that are related to the ionization state of the wind (\\S5). ",
        "conclusions": "%  Our results indicate that X-ray flux appears to be related to the state of the wind (as observed in the H$\\alpha$ emission strength). When the H$\\alpha$ emission is strong the X-ray flux is consistently low, but when the emission is relatively weak the X-ray flux can cover a wide range in values.   There are three processes that can potentially explain these observations: (1) X-ray photoionization of the wind leading to a decrease in H$\\alpha$ emissivity; (2) wind-related changes in column density towards the X-ray source; and (3) variations in wind strength that result in changes in accretion rate.  Here we argue that all three processes are required to help explain our results and that changes in the wind accretion rate may trigger the X-ray transitions in Cyg~X-1.  The first possible explanation is that increased X-ray emission leads to heating of the wind and a corresponding decrease in H$\\alpha$ emissivity \\citep{bro99b}.  \\citet{vlo01} show how the X-ray ionization of the wind in the low/hard state causes orbital variations in the ultraviolet P~Cygni lines formed in the wind (e.g., the Hatchett-McCray effect; \\citet{hat77}).  The X-ray ionization effects become more severe in the high/soft state, and \\citet{wen99} describe how the wind absorption effects decrease significantly in such highly ionized conditions. Hydrogen is expected to be nearly fully ionized in the base of the wind irrespective of the X-ray state, but if the gas is heated by an elevated soft X-ray flux, then the H$\\alpha$ emissivity will drop ($\\propto T^{-1.2}$; \\citet{ric98}).  The best evidence for ionization-related variability is the reduction in H$\\alpha$ equivalent width we observed during the mini-flare of X-ray flux around HJD 2,451,426.8 (Fig.~11). The increase in soft X-ray flux and reduction in H$\\alpha$ occurred more-or-less simultaneously (within the time resolution of the data), as expected for the small light-travel time between the X-ray source and the facing hemisphere of the supergiant ($\\approx 54$ s). On the other hand, the characteristic wind crossing time between components is approximately 10 hours, and our data, although sparse in time coverage, does not indicate any time lag of this order between the H$\\alpha$ and X-ray variations.   This event occurred near orbital phase $\\phi=0.77$, and the two spectra we obtained then show H$\\alpha$ profiles with unusually deep absorption cores (see Fig.~4) compared to observations on the preceding and following nights.  At this orientation the wind between the stars was moving tangentially to our line of sight and the emission from this region, if present, would have filled in the line core.  However, if this inner region was significantly photoionized by the X-ray mini-flare, then the residual emission in the line core would have vanished to produce the deeper absorption core.  Thus, the timing and line properties of this event suggest it was caused by X-ray photoionization.  On the other hand, there are several lines of evidence that X-ray photoionization is not the dominant cause of the H$\\alpha$ variations.  First, we would expect the anti-correlation between H$\\alpha$ emission and X-ray flux to be better defined than exhibited in Figure~13 if X-ray photoionization drove the emission variations (although the scatter may be partially the result of the poor time resolution and poor overlap of the H$\\alpha$ and X-ray observations).  There are several examples (especially during the X-ray low/hard state) where we observed large excursions in the H$\\alpha$ emission while the X-ray flux was essentially constant (Fig.~10). Second, if X-ray photoionization significantly altered the H$\\alpha$ emissivity of the gas above the hemisphere facing the X-ray source, then we should observe changes in the shape of the supergiant's P~Cygni component with orbital phase.  The variations would be largest in the high/soft X-ray state and would result in decreased blue absorption near $\\phi = 0.5$ and decreased red emission near $\\phi = 0.0$. The spectra observed in the high/soft X-ray state (Fig.~5) show that the predicted changes may be present but are relatively minor in nature.   Third, we find evidence of large scale variations in the H$\\alpha$ emission that forms in the X-ray shadow above the hemisphere facing away from the X-ray source where no X-ray photoionization should occur. This region in the wind is best isolated in the radial velocity distribution of the emission near orbital phase $\\phi=0.5$ when this backside outflow is the dominant contributor to the red-shifted peak of the P~Cygni profile.   However, given the probable low orbital inclination, this part of the profile also includes some contributions from the X-ray illuminated hemisphere.  We show in Figure~15 the red peak emission height for spectra obtained near this phase plotted against X-ray flux.   These diagrams show that the strength of the emission component from the X-ray shadow region is generally independent of the X-ray flux level in both the low/hard ({\\it plus signs}) and high/soft ({\\it diamonds}) X-ray states, as expected for an emitting region sheltered by the supergiant. In contrast, we show in Figure~16 the peak intensity of the blue emission peak isolated near orbital phase $\\phi = 0.25$ that displays evidence of an anti-correlation between H$\\alpha$ emission and X-ray flux as expected for irradiated gas close to the black hole (\\S3). The large range in the strength of the H$\\alpha$ emission from the X-ray shadow region suggests that variations are caused by factors other than changes in X-ray photoionization. Since the emission strength is very sensitive to gas density ($\\propto n^2$), we conclude that the main source of the H$\\alpha$ emission variations is the fluctuation in basal wind density rather than ionization state.  \\placefigure{fig15}     %  \\placefigure{fig16}     %  If the H$\\alpha$ variations primarily reflect changes in wind density close to the supergiant, then it is possible that the X-ray flux is partially modulated by the changing absorption in the stellar wind (particularly important for soft X-rays). We doubt that wind absorption plays a major role in explaining the general anti-correlation between H$\\alpha$ and X-ray emission because the scatter in their temporal variations (Fig.~13) is much larger than we would expect for a direct cause-and-effect relationship. Furthermore, the amplitude of the wind density fluctuations implied by the H$\\alpha$ changes is too small to explain the range in X-ray variability. \\citet{wen99} show that the X-ray light curve can be explained by the changing column density of the line of sight to the black hole as we peer through different portions of the supergiant's wind (strongest absorption at supergiant inferior conjunction, $\\phi=0.0$). Their ASM 1.5 -- 3 keV orbital light curve for the low/hard state shows a 25\\% decrease at $\\phi=0.0$ relative to $\\phi=0.5$.  \\citet{bal00} also studied the variation in column density with orbital phase, and they suggest that the orbital modulation in X-ray flux corresponds to a fluctuation of $6 - 10\\times$ in column density (depending on the assumed ionization of the wind). However, the observed variations in the emission strength generally indicate column density changes of a factor of $<2$ (see below).   Wind column density changes of this order are too small to account for the large changes in X-ray flux. Thus, wind absorption variations are insufficient to explain the general anti-correlation between the H$\\alpha$ emission strength and X-ray flux. Nevertheless, we might expect the wind absorption processes to appear most prominent near inferior conjunction of the supergiant when the X-ray light curve attains a minimum and a peak occurs in the frequency of rapid X-ray dips \\citep{wen99,bal00,fen02}.  It is noteworthy in this regard that the case of the simultaneous H$\\alpha$ maximum and X-ray flux minimum shown in Figure~12 did indeed occur near this phase (at $\\phi = 0.02$).  If we ignore the effects of X-ray ionization on the H$\\alpha$ emissivity, then we can make an approximate estimate of how the wind mass loss rate and density vary as a function of the H$\\alpha$ emission strength \\citep{pul96}. \\citet{her95} show predictions of how the H$\\alpha$ profile will vary as a function of mass loss rate (see their Fig.~4), and we measured the equivalent widths of their profiles to calibrate the mass loss rate as a function of H$\\alpha$ equivalent width (prorated to their final estimate of mass loss rate for the time of their observations, $3.0\\times 10^{-6}$ $M_\\odot$~y$^{-1}$).   The functional fit in this case is \\begin{equation} -\\dot{M} = (1.85 -1.01~W_\\lambda - 0.04~W_\\lambda^2)~\\times 10^{-6}~M_\\odot~{\\rm y}^{-1}. \\end{equation} We grouped our equivalent width data according to the X-ray state at the time of observation (setting aside the results obtained near HJD 2,451,896 that may correspond to a ``failed transition''), and the mean equivalent width for each group yields mass loss rates of $(2.57\\pm 0.05)\\times 10^{-6}$ and $(2.00\\pm 0.03)\\times 10^{-6}$ $M_\\odot$~y$^{-1}$ for the low/hard and high/soft states, respectively (the quoted errors are based on the standard deviation of the mean and do not include the larger errors associated with the calibration of the $W_\\lambda - \\dot{M}$ relationship and with our neglect of the multiple component nature of the H$\\alpha$ emission). This suggests that the mass loss rate is $\\approx 22\\%$ lower during the rare high/soft state compared to the more common low/hard state. Since this calculation ignores the X-ray photoionization that may be more important in the high/soft state and since photoionization will decrease H$\\alpha$ emission strength, our estimate for the mass loss rate during the high/soft state is best regarded as a lower limit. Nevertheless, the general decrease in wind strength during the high/soft state is also observed in the emission from the X-ray shadow region (Fig.~15) that should be relatively free from the effects of photoionization.  We find that the mean residual emission intensity of the red peak near phase $\\phi = 0.5$ is $0.139\\pm 0.008$ and $0.114\\pm 0.006$ for the low/hard and high/soft states, respectively. \\citet{wen99} also found that the wind mass loss rate is lower in the high/soft state based upon their analysis of the X-ray orbital light curve.  Taken at face value, our results present a paradox: the X-ray flux decreases when the wind mass loss rate increases.    We suggest that the resolution of this quandary lies in how the wind velocity changes with wind ionization state (originally proposed by \\citet{ho87} and developed in more physical terms by \\citet{ste91}). When the wind mass loss rate is high, the wind density is also proportionally high, and therefore the X-ray ionization effects on the wind are confined to the region close to the black hole since the size of the surrounding ionized region depends on wind density as $n^{-1}$ \\citep{ste91,vlo01}. Most of the wind volume surrounding the star will contain the many important ions that propel the radiative driving of the wind, so that it reaches a terminal velocity of approximately 2100 km~s$^{-1}$ \\citep{her95}.   However, the wind flow close to the black hole will become ionized, and these advanced ionization states will generally have transitions at frequencies much higher than the peak of the stellar flux distribution.  Consequently there are fewer absorbing transitions at frequencies where the stellar flux is concentrated and radiative driving of the wind becomes less effective.  Stellar wind gas leaving the part of the supergiant facing the X-ray source will be accelerated by radiative driving, reaching a significant fraction of the terminal velocity before crossing the distant ionization boundary \\citep{ste91}.  Models of wind accretion by the black hole suggest that the mass accretion rate is proportional to $\\dot{M}/ v^4$ \\citep{bon44,lam76,ste91}, and since the flow velocity $v$ is relatively large, only modest mass accretion occurs.   It is easy to imagine that some of the outflow from the supergiant to the black hole bypasses the black hole altogether, and absorption by the gas beyond the black hole could explain the presence of a secondary minimum in the low/hard state X-ray light curve at orbital phase $\\phi=0.5$ \\citep{kar01}.  However, when the mass loss rate drops, the X-ray ionization zone will become larger because of the reduced wind density. The dominant ions that provide the resonant transitions that in turn drive the stellar wind will disappear with increased ionization, and the wind outflow will experience only a modest acceleration and reach a speed of just a few hundred km~s$^{-1}$ (before the flow dynamics become dominated by the gravitational acceleration of the black hole). This altered outflow will be denser and slower in the vicinity of the black hole, and since the accretion rate varies as $\\dot{M}/ v^4$, the overall accretion rate will increase significantly because of the slower outflow (and despite a real decline in $\\dot{M}$).  If this basic scenario is correct, then transitions from the low/hard to the high/soft state are triggered when the supergiant undergoes an episode of reduced mass loss.   \\citet{ste91} demonstrates that the decrease in the wind force multiplier occurs rather suddenly once a specific gas density is reached, and we would argue that this is the reason why the transitions are relatively fast and the X-ray states are bimodal. Once the transition occurs, the system tends to remain in the high/soft state because the increased accretion rate and associated larger X-ray fluxes make it easier to keep the wind in the highly ionized state (even with modest increases in mass loss rate). The line of sight to the black hole then passes through mainly ionized gas all around the orbit, and the wind absorption of the X-ray flux decreases so much that the X-ray light curve modulation vanishes \\citep{wen99}. The increased mass accretion in this state causes the optically thick, geometrically thin (Keplerian) accretion disk to extend further inwards towards the black hole, producing more soft X-rays, while the coronal (ADAF or sub-Keplerian) region, which produces the bulk of the hard X-rays, becomes smaller, and thus the X-ray spectrum softens \\citep{ebi96,esi98,bro99b}. The enhanced soft X-ray flux heats the accreting gas, which reduces the H$\\alpha$ emission from the accretion flow (Fig.~9, 16). It is only once the supergiant can maintain a strong mass outflow that the wind becomes dense enough to reduce the ionization effects and to create the high speeds which then lower the mass accretion rate and force the system back to the low/hard state.   The long term variations in the mass loss rates of massive supergiant stars are not well documented, but there is circumstantial evidence of significant variations on time scales of years \\citep{ebb82,mar02}.   Thus, we suggest that the high/soft states that occur every 5 years or so in Cyg~X-1 correspond to quasi-cyclic minima in the mass loss rate of the supergiant."
    },
    "0206/astro-ph0206123_arXiv.txt": {
        "abstract": "We have observed the nearby spiral galaxy NGC6503 using Chandra. Seven discrete sources associated with the galaxy have been found, one of them coincident with its Liner-starburst nucleus. One of the sources corresponds to a ULX with L$_{\\rm x} \\ga 10^{39}$ ergs s$^{-1}$. Previous ROSAT observations of the galaxy show that this source has varied by at least a factor 3 in the last 6 years. No evidence is found for strong diffuse emission in the nuclear region or the presence of a low luminosity AGN. ",
        "introduction": "NGC6503 is a low luminosity ($M_{B} = -18.28$) edge-on Sc galaxy at a distance of 7Mpc. It has been the subject of several kinematic studies which show an unexpected drop in the stellar velocity dispersion in the inner region of the galaxy (Bottema \\& Gerritsen 1997). Based on the large observed [NII]/H$\\alpha$ and [SII]/H$\\alpha$ line ratios Ho, Filippenko \\& Sargent (1995) classified the nucleus as a transition-Seyfert-2 nucleus. With a narrow H$\\alpha$ luminosity $\\sim 4\\times 10^{37}$ ergs s$^{-1}$ this would be one of the lowest luminosity Seyfert nucleus known.  However, diagnostic diagrams using the [OI]$\\lambda 6300$, [OII]$\\lambda 3717$ and [OIII]$\\lambda 5007$ optical lines has shown that the nuclear activity in NGC6503 is better classified as borderline between starburst and LINER (Lira etal 2002) and therefore the presence of a low luminosity AGN has been put into dispute.  ROSAT HRI X-ray observations of NGC6503 were reported by Lira etal (2000). The image of the galaxy showed an extended nuclear source of luminosity $\\sim 10^{39}$ ergs s$^{-1}$ in the 0.5-2.4 keV energy range. The presence of this source, together with the early Seyfert-like classification of the nucleus, made this galaxy a interesting target for high resolution Chandra observations, which would enable us to study the low luminosity active nucleus and the circumnuclear region in detail.  In this paper, we show that the X-ray Chandra observations confirm that NGC6503 is an example of the activity seen in normal spiral galaxies harbouring a moderate starburst nucleus, while no evidence for the presence of an AGN is found.   \\begin{figure*} \\centering \\begin{minipage}[c]{0.7\\textwidth} \\centering \\begin{Huge} n6503.jpg \\end{Huge} \\end{minipage}% \\begin{minipage}[c]{0.3\\textwidth} \\centering \\includegraphics[angle=90,scale=0.25]{Nuc_reg.ps} \\includegraphics[scale=0.25]{C.ps} \\includegraphics[scale=0.25]{D.ps} \\end{minipage} \\caption{Left: JKT grey-scale R band image of NGC6503 with the position of the detected discrete sources marked with circles. Right: Detail of (from top to bottom) the nuclear region, source C and source D, with overlaid X-ray contours. The grey-scale images correspond to HST WFPC2 observations (PIs: Davis, Illingworth) obtained with the filters F606W (for the nuclear image) and F814W (for the background objects - see section 5.3).} \\end{figure*}  \\begin{table*} \\caption{Detected discrete sources in NGC6503. Background subtracted source counts ($SC$) in the 0.3-10 keV range are given for each source. The significance of the detections was determined as $SC/\\sigma_{B}$, where $\\sigma_{B}$ is the standard deviation of the background and was computed as $1 + \\sqrt{\\rm{Background\\ counts} + 0.75}$. This correction to the standard deviation gives a more appropriate estimate of Poissonian errors for cases of low number counts (Gehrels 1986). {\\it Observed\\/} fluxes are in units of $10^{-14}$ ergs s$^{-1}$ cm$^{-2}$. {\\it Intrinsic\\/} luminosities and fluxes (for those sources not associated with NGC6503) are in units of ergs s$^{-1}$ and ergs s$^{-1}$ cm$^{-2}$, respectively. Hydrogen column densities are in units of $10^{21}$ cm$^{-2}$. $a$: Results from spectral fitting for source 6 (see text).}  \\centering \\begin{tabular}{lllcrccccl} \\hline Source & \\multicolumn{2}{c}{Position (J2000)} & Counts & Signif. & $\\Gamma$ & $N_{H}$ & F$^{\\rm obs}_{\\rm x}$ & L$^{\\rm int}_{\\rm x}$--F$^{\\rm int}_{\\rm x}$ & Comments\\\\ \\hline 1 & 17 49 12.48  &70 9 31.4&  $136 \\pm 17$   \t&65.2\t&3.0\t&1.0\t&4.69\t&$3.35\\times10^{+38}$ &\\\\ 2 & 17 49 26.43  &70 8 39.7&  $13  \\pm 5$ \t&6.0\t&1.5\t&1.0\t&0.90\t&$5.85\\times10^{+37}$ & NGC6503 Nucleus\\\\ 3 & 17 49 27.89  &70 8 36.7&  $34  \\pm 7$ \t&16.1\t&1.5\t&1.0\t&2.34\t&$1.52\\times10^{+38}$ &\\\\ 4 & 17 49 28.84  &70 8 32.7&  $34  \\pm 7$ \t&16.1\t&1.5\t&1.0\t&2.34\t&$1.52\\times10^{+38}$ & Opt ID: GC in NGC6503\\\\ 5 & 17 49 29.09  &70 8 44.2&  $154 \\pm 13$   \t&73.8\t&1.5\t&1.0\t&10.6\t&$6.89\\times10^{+38}$ &\\\\ 6 & 17 49 31.72  &70 8 20.3&  $255 \\pm 17$   \t&122.4\t&1.5\t&5.0\t&25.8\t&$1.98\\times10^{+39}$ & ULX in NGC6503\\\\ 6$^{a}$&         &         &                 \t&\t&$kT=4$ keV&5.5\t&17.4\t&$1.46\\times10^{+39}$ &\\\\ 7 & 17 49 38.09  &70 8 31.9&  $7   \\pm 4$ \t&3.2\t&1.5\t&1.0\t&0.48\t&$3.12\\times10^{+37}$ &\\\\ A & 17 49 09.84  &70 7 26.4&  $9   \\pm 4$ \t&4.1\t&2.0\t&0.4\t&0.62\t&$6.97\\times10^{-15}$ & Opt ID: Bkg QSO\\\\ B & 17 49 11.26  &70 6 24.1&  $8   \\pm 4$ \t&3.6\t&2.0\t&0.4\t&0.55\t&$6.19\\times10^{-15}$ & Opt ID: Bkg QSO\\\\ C & 17 49 18.58  &70 6 31.8&  $24  \\pm 6$ \t&11.3\t&2.0\t&0.4\t&0.83\t&$1.03\\times10^{-14}$ & Opt ID: Bkg QSO\\\\ D & 17 49 26.65  &70 6 51.8&  $31  \\pm 7$ \t&14.7\t&2.0\t&0.4\t&2.14\t&$2.41\\times10^{-14}$ & Opt ID: Bkg QSO\\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "We have obtained high resolution Chandra X-ray observations of the nearby spiral galaxy NGC6503. The data have shown the presence of faint diffuse emission around the central region of the galaxy with a luminosity of $\\sim 10^{39}$ ergs s$^{-1}$ in the 0.3-10 keV energy range. A total of 11 compact sources were identified in the field of view, 7 of which seem to be associated with NGC6503. The remaining sources seem to correspond to background quasars.  A weak, L$_{\\rm x} \\sim 6\\times10^{37}$ ergs s$^{-1}$, source is found to be coincident with the Liner-starburst nucleus of the galaxy. Another X-ray source has a faint optical counterpart which might correspond to a globular cluster in NGC6503. Finally, a L$_{\\rm x} \\ga 10^{39}$ ergs s$^{-1}$ ULX has been identified which is nearly coincident with a distinctive knot of blue emission seen in the disk of the galaxy. Previous ROSAT observations obtained 6 years before show that this source has varied by at least a factor 3."
    },
    "0206/astro-ph0206315_arXiv.txt": {
        "abstract": "In this work we present a model of the universe in which dark energy is modelled explicitely with both a dynamical quintessence field and a cosmological constant. Our results confirm the possibility of a future collapsing universe (for a given region of the parameter space), which is necessary for a consistent formulation of string theory and quantum field theory. We have also reproduced the measurements of modulus distance from supernovae with good accuracy. ",
        "introduction": "From 1998 to date several important discoveries in the astrophysical sciences have being made, which have given rise to the so called New Cosmology \\cite{turner1,turner2}. Amongst its more important facts we may cite: the universe expands in an acelerated way \\cite{riess,perlmutter}; the first Doppler peak in the cosmic microwave background is strongly consistent with a flat universe whose density is the critical one \\cite{cmb}, while several independent observations indicate that matter energy density is about one third of the aforementioned critical density \\cite{Om,smoot}. The last two facts implied that some unknown component of the Universe ''was missing'', it was called dark energy, and represents near two thirds of the energy density of the universe. The leading candidates to be identified with dark energy involve fundamental physics and include a cosmological constant (vacuum energy), a rolling scalar field (quintessence), and a network of light, frustrated topological defects \\cite{mst}.  On the other hand, an eternally accelerating universe seems to be at odds with string theory, because of the impossibility of formulating the S-matrix. In a de Sitter space the presence of an event horizon, signifying causally disconnected regions of space, implies the absence of asymptotic particle states which are needed to define transition amplitudes \\cite {banks,cline}. This objection against accelerated expansion also applies to quantum field theory (QFT)\\cite{sasaki}.  Due to the above there is a renewed interest in exponential quintessence, because in several scenarios exponential potentials can reproduce the present acceleration and predict future deceleration, so again string theory has well defined asymptotic states \\cite{cline,kl}. Worthwhile to notice that exponential quintessence had been so far overlooked on fine tuning arguments, but several authors have recently pointed out that the degree of fine tuning needed in these scenarios is no more than in others usually accepted \\cite{cline,kl,rubano}.  The cosmological constant can be incorporated into the quintessence potential as a constant which shifts the potential value, especially, the value of the minimum of the potential, where the quintessence field rolls towards. Conversely, the height of the minimum of the potential can also be regarded as a part of the cosmological constant. Usually, for separating them, the possible nonzero height of the minimum of the potential is incorporated into the cosmological constant and then set to be zero. The cosmological constant can be provided by various kinds of matter, such as the vacuum energy of quantum fields and the potential energy of classical fields and may also be originated in the intrinsic geometry. So far there is no sufficient reason to set the cosmological constant (or the height of the minimum of the quintessence potential) to be zero \\cite{hwang}. In particular, some mechanisms to generate a negative cosmological constant have been pointed out \\cite{ss,gh}.  The goal of this paper is to present a model of the universe in which the dark energy component is accounted for by both a quintessence field and a negative cosmological constant. The quintessence field accounts for the present stage of accelerated expansion of the universe. Meanwhile, the inclusion of a negative cosmological constant warrants that the present stage of accelerated expansion will be, eventually, followed by a period of collapse into a final cosmological singularity (AdS universe). ",
        "conclusions": "In a recent paper \\cite{hwang} it is pointed out that the ultimate fate of the evolution of our universe is much more sensitive to the presence of the cosmological constant than any other matter content. In particular, the universe with a negative cosmological constant will always collapse eventually, even though the cosmological constant may be nearly zero and undetectable at all at the present time. Our results support the very general assertions of \\cite{hwang}, we have shown that for a determined region of the parameter space, the universe collapses. This also favours the consistent formulation of string theory and quantum field theory, as explained in the introduction. The experimental measurements of modulus distance from the supernovae are adequately reproduced within an accuracy of 0.5$\\%$. So far, we have investigated one of the several possible branches of the solution, leaving for the future the investigation of the others. We have also reserved for future work the careful examination of this universe near its beginning (i.e., just after the decoupling of matter and radiation).   We acknowledge Claudio Rubano, Mauro Sereno and Paolo Scudellaro, from Universita di Napoli \"Federico II\" , Italy, for useful comments and discussions and Andro Gonzales for help in the computations."
    },
    "0206/hep-th0206044_arXiv.txt": {
        "abstract": " ",
        "introduction": "Recently, there has been some interest in multigravity theories \\cite{Kogan:1999wc,Gregory:2000jc,Kogan:2000cv,Kogan:2000xc,Kogan:2000vb,Kogan:2001yr}  where gravity is modified at cosmological scales. These theories involve brane configurations in a higher than four dimensional spacetime where normal four dimensional  gravity on the branes is modified  in the far infrared due to the presence of a massive (but ultralight) graviton component in the low energy theory (see \\cite{Kogan:2001ub} for a review and \\cite{Papazoglou:2001cc} for a detailed presentation).   The attractive feature of these theories is that they provide an alternative observational window to extra dimensional physics which is testable in current observations \\cite{Will:1997bb,Binetruy:2000xv,Uzan:2000mz,Bastero-Gil:2001rv,Choudhury:2002pu}. Most importantly, these modifications of gravity are at such scales from which current observations indicate a dark energy component in our universe \\cite{Riess:1998cb,Perlmutter:1998np}. It would be tempting to attribute this to the dynamics of a multigravity system.   The basic idea in constructing multigravity models is to localize gravity at the same time in different places along the extra dimension(s). Once one has a higher dimensional brane configuration which localizes gravity, {\\textit{e.g.}} \\cite{Randall:1999vf,Gherghetta:2000qi}, the low energy effective theory is governed by a massless graviton field. By multilocalizing gravity \\cite{Kogan:2001wp} in a superposition of such configurations, the degeneracy of the massless modes will be lifted and the low energy theory will contain apart of a massless mode, a collection of light massive gravitons. How light these gravitons will be depends on how strong the localization in the single graviton  configuration is. In the particular models that have been examined in the literature \\cite{Kogan:1999wc,Gregory:2000jc,Kogan:2000cv,Kogan:2000xc,Kogan:2000vb,Kogan:2001yr}, the localization was exponentially strong  and thus the mass splittings between  the light gravitons and the remaining of the Kaluza-Klein (KK) spectrum was exponentially large. This gave the opportunity to realize models which escaped observational bounds and had  interesting phenomenological implications. An alternative mechanism which has similar effects arises in the case where four dimensional gravity is induced on the brane due to quantum loops of  matter living on the brane \\cite{Dvali:2000hr,Dvali:2000xg}. In that case, the old result of Sakharov \\cite{Sakharov:pk} (for a review of induced gravity see \\cite{Adler:1982ri}) was exploited to render the geometry on a brane, embedded in a flat higher dimensional space, four dimensional.  However, so far the models \\cite{Kogan:1999wc,Gregory:2000jc,Kogan:2000cv,Kogan:2000xc,Kogan:2000vb,Kogan:2001yr}  where studied at the linearized level which becomes invalid when speaking about cosmological distance dynamics. One clearly has to go beyond the linear theory, which is the main aim of this paper. Nonlinear bigravity theories were first introduced in the seventies as effective descriptions of a sector of hadronic physics \\cite{Isham:gm}. It is argued in a companion paper \\cite{DK} that nonlinear bigravity theories can arise in several different (purely gravitational) contexts: multibrane configurations, certain classes of Kaluza-Klein models, some types of non-commutative geometry models, {\\textit{etc.}} It is, therefore, important to try to delineate what are the generic predictions of classes of bigravity theories. One of the main conclusions of the present paper is that bigravity naturally gives rise to a late period of cosmic acceleration. Bigravity can then be used as a new theoretical model of dark energy (with specific anisotropic features, in certain cases, that make it phenomenologically distinguishable from quintessence models).   Accelerating solutions were also found in the context of one particular model of brane-induced gravity \\cite{Dvali:2000hr,Dvali:2000xg}, and their phenomenological consequences have been explored in detail as a possible theoretical model of dark energy \\cite{Deffayet:2000uy,Deffayet:2001pu}. Both models share the common feature of modifications of gravity at large scales. However, our work differs in several ways. We study general classes of four-dimensional effective theories instead of one particular five-dimensional model. We study general classes of solutions of these theories including their stability. We find several different types of accelerating solutions some being isotropic, but many featuring a novel type of anisotropic acceleration.    In the present paper, we  will firstly describe the formalism needed for  discussing a general bigravity system. We will be interested in discussing cosmological solutions for such a system. As illustrative models  we will use two types of potentials coupling the two metrics. One which resembles the Pauli-Fierz mass term in the linearized theory and one which is motivated by a higher dimensional brane construction. For these potentials we will be examining the most simple case, by imposing special symmetries, and without including matter. We introduce a description of the coupled cosmological evolution of the two metrics in terms of a ``mechanical'' model: two ``relativistic particles'' connected by a nonlinear ``spring''.  We will see that there is a generic period of acceleration of one or both of the metrics. For the Pauli-Fierz potential we discover two classes of accelerating solutions. One where anisotropies play a crucial role, and one where the cosmology is isotropic. On the other hand, for the brane potential our symmetry requirements only allow for solutions where anisotropies play an important role. In the final state the relative lapse between the two metrics tends to run away to infinity, when using the above illustrative coupling potentials. This run-away signals the breakdown of our effective theory. We discuss more general classes of potentials which naturally lead to a confinement of the relative lapse within a limited range. Such models lead to an interesting ``locking'' mechanism of the evolution of the two metrics. At the end, we briefly discuss the inclusion of matter in the above systems and propose the multigravity scenario as a candidate for a purely gravitational type of dark energy.  In the following we adopt the mostly plus metric signature $(-,+,+,+)$ and use the following definition for the Riemann tensor  $R^K_{~\\Lambda MN}=\\de_M \\Gamma^K_{\\Lambda N}-\\de_N \\Gamma^K_{\\Lambda M}+\\Gamma^H_{\\Lambda N}\\Gamma^K_{MH}-\\Gamma^H_{\\Lambda M}\\Gamma^K_{NH}$. We use  capital letters to label four dimensional spacetime coordinates and lower case letters to label three dimensional space coordinates. For the coordinate basis we use greek letters $M,N,\\Lambda,\\dots=0,1,2,3$ and $\\mu,\\nu,\\lambda,\\dots=1,2,3$,  and for the vierbien latin ones $A,B,C,\\dots=0,1,2,3$ and  $a,b,c,\\dots=1,2,3$. The reason for using, somewhat unconventionally, $\\mu,\\nu,\\lambda,\\dots$ for spatial indices will be explained below. ",
        "conclusions": "We have explored general spatially-flat cosmological solutions (of the anisotropic Bianchi I type) of classes of nonlinear bigravity theories. Even within this restricted class of homogeneous cosmologies we focused on special cases. We did not explore the possibility (mentioned in Section 3) where, due to a type of spontaneous symmetry-breaking mechanism, the relative shift vector $b^{\\mu} \\equiv e^{-\\bar{\\gamma}}(b_2^{\\mu}-b_1^{\\mu})$ be non-zero. After a brief discussion of the structure of the evolution system for two coupled spatial metrics, we restricted ourselves to the simple case where the two metric tensors can be simultaneously diagonalized. Even this simplified case leads to very rich dynamics which can be conveniently described in terms of a mechanical model (represented in Fig.\\ref{genplot}): two ``relativistic particles'', moving in a $(2+1)$-dimensional Lorentzian space, and connected by some non-linear ``spring'', {\\textit{i.e.}} interacting via some bigravity potential ${\\cal V}_{12}=(g_1g_2)^{1/4}V({\\bf g}_1^{-1}{\\bf g}_2)$.  One of the first important conclusion of our study is that the long-term behaviour of this coupled system crucially depends on the ability of the potential $V$ to confine the evolution of the relative lapse $\\gamma=\\log\\left(N_2/N_1\\right)$. Due to the former ``gauge nature''  of $\\gamma$ (when the potential is absent), {\\textit{i.e.}} the absence of kinetic terms for $\\gamma$, the equation of motion of $\\gamma$ is algebraic. We found that the continued existence (in the long term)  of a solution for $\\gamma$ sensitively depends on the nature of the function $V(\\gamma)$. For instance, we found that the potential $V({\\bf g}_1^{-1}{\\bf g}_2)$ derived from the five-dimensional brane constructions has only a marginal ability to confine the evolution of $\\gamma$ to a limited range of variation. Indeed, we found that many (and maybe most) solutions of bi-cosmology, with such a brane potential, evolve, after some finite time, into a state where $\\gamma$ quickly runs away towards infinity. By exploring the behaviour of physical observables near the moment where $|\\gamma| \\to \\infty$, we have shown that the run away does not correspond to any observable singularity in either of the two metrics. We have argued that this run away only signals a breakdown of the effective four-dimensional description that we use. Indeed, it seems that, as $|\\gamma| \\to \\infty$, some previously heavy modes become light and should now be taken into account in the effective action.  Having understood the root of this run away behaviour, we have focused our physical study of bigravity on the class of potentials $V(\\gamma)$ which have the strong-enough confining property with respect to $\\gamma$. A simple example of a potential in this class is the ``quadratic plus quartic'' Pauli-Fierz-type potential (\\ref{defpot}). Such a potential allows for solutions which evolve on long-time scales, without encountering any breakdown linked to a run away of $\\gamma$. [Note, however, that we do not claim that this is true for all solutions. It is certainly possible to concoct initial data leading to a $\\gamma$ runaway after a finite time.]  When using such $\\gamma$-confining potentials (or when considering, as we do in most of the text, the effect of any potential up to times smaller than the moment of quick $\\gamma$ runaway) we have found that the qualitative behaviour of generic cosmological solutions can be nicely understood in terms of the mechanical model of Fig.\\ref{genplot}. For instance, when the separating vector $\\delta^{\\mu}=\\beta^{\\mu}-\\alpha^{\\mu}$ between the two ``particles'' is spacelike, and the potential is polynomial in $\\delta^{\\mu}$ (and attractive, as the modified Pauli-Fierz potential (\\ref{defpot})), the coupled motion of $\\alpha^{\\mu}$ and $\\beta^{\\mu}$ is similar to slow-roll inflation. The separating vector $\\delta^{\\mu}$ plays the role of the inflaton, and drives an exponential-type expansion of the vertical position of the ``center of mass'' ${1 \\over 2}\\sigma={1 \\over 2}(\\beta+\\alpha)$ (which represents the average volume of the two metrics). A qualitatively new feature of this type of bigravity slow-roll inflation is its growing anisotropy. As the connecting vector gets smaller, each metric tends to expand more and more differently in three spatial directions (linked to the ``direction'' of the vector $\\delta^{\\mu}$). This anisotropic slow-roll inflation ends up in a regime where the ``spring'' connecting the two ``particles'' makes them oscillate along a spacelike direction (see Fig.\\ref{spaceplotnum}). Similarly to the oscillatory period following slow-roll (for a chaotic inflation type potential, {\\textit{e.g.}} $V(\\phi)={1 \\over 2}m^2\\phi^2$), these bigravity oscillations lead to a power-law expansion law. We expect such anisotropic accelerating solutions to exist in a general multigravity theory, and thus also in the brane-induced model   \\cite{Dvali:2000hr,Dvali:2000xg}, which, as we said above, can be viewed as the $N \\to \\infty$ limit of the particular ``nearest neighbour'' interactions multigravity model.   When the separating vector $\\delta^{\\mu}=\\beta^{\\mu}-\\alpha^{\\mu}$ is timelike, we found (when using the $\\gamma$-confining potential (\\ref{defpot})) a remarkable phenomenon of ``locking'' of the two metrics. [We have indicated in Eqs.(\\ref{dgcond}), (\\ref{dcond}), (\\ref{lockedrates}) the general conditions under which this ``locking'' phenomenon occurs.] In the visual language of Fig.{\\ref{genplot}}, the two ``particles'' lock in a perpetual ``chasing'' configuration where their vertical separation tends to a non-zero constant, while their ``center of mass''  continues to move upwards. In bimetric language, the locking corresponds to a bi-de-Sitter configuration: each metric expands exponentially, and the two expansion rates are equal (in the average proper time). Contrary to the spacelike case where anisotropies played an important role, here this configuration is obtained for isotropically expanding metrics. It would be interesting to explore the basin of attraction of such a locked state among generic (timelike-type)  bigravity evolutions. As indicated above, these locked configurations admit (provided some system of $N-1$ equations for $N-1$ unknowns admit a real solution) a multi-de-Sitter generalization in a general multigravity model. [As we mentioned, this is a way of interpreting the solution of   \\cite{Deffayet:2000uy,Deffayet:2001pu}.] Since we found that the bigravity locked solution was stable, we expect this feature to extend to the multigravity case.  From the phenomenological point of view, one of the major conclusions of this work is that bigravity cosmologies generically exhibit a period of cosmic acceleration for one or both of the metrics. This conclusion applies even to the case of ``bad'' potentials which cannot permanently constrain the evolution of $\\gamma$ to a bounded range. This result suggests that bigravity could be the origin of the observed cosmic acceleration, {\\textit{i.e.}} that it could be the the source of {\\textit{dark energy}}. In other words bigravity naturally defines a kind of ``tensor quintessence''. In brane models, the mass parameter $m^2$ appearing in the potential $V$ is an exponentially decreasing function of the interbrane distance. It is therefore not unnatural to have an $m^2$ as small as it is required to explain the observed cosmic acceleration ({\\textit{i.e.}} $m \\sim 10^{-33}$eV)\\footnote{Let us note that the parameter $\\mu$ appearing in (\\ref{a1}) would then be $\\mu \\sim 10^{-3}$eV.}.  Our preliminary studies of the transition between matter domination and vacuum domination seem to indicate that (at least for classes of potentials) this transition can be as smooth as in the usually considered dark-energy models (such as a cosmological constant, or some type of scalar quintessence). It is, however, interesting to note that, at least in the spacelike separated case, bigravity makes qualitatively new predictions: it predicts a growing anisotropy of the expansion of the universe. It would be interesting to study the imprint of this phenomenon (which started to take place only ``recently'', {\\textit{i.e.}} for redshifts $z \\lesssim 0.5$) on observable phenomena, and notably on the Cosmic Microwave Background.   Finally,   on a more speculative view, it would be interesting to explore the possibility that bigravity explains the primordial inflation needed to explain the gross features of our universe. For this, one would probably need a mass scale of order $m \\sim 10^{-6} M_{\\rm Pl}$. If we contemplate a ``spacelike'' scenario, the needed large initial value of $\\delta^{\\mu}$ to have a long stage of bigravity slow roll inflation, might be naturally provided by the recently discovered generic chaotic behaviour taking place in (bulk) string/M cosmology \\cite{Damour:2000wm,Damour:2000hv,Damour:2001sa}. Indeed, the chaotic behaviour naturally leads, near $t \\sim t_{\\rm string}$ to very large ``oscillations'' in the (logarithmic) scale factors $\\alpha^{\\mu}$ of the metric (considered at some spatial points). When comparing two metrics (either at two different bulk points, or on two branes) it is then natural to reach large values of $\\beta^{\\mu}-\\alpha^{\\mu}$. On the other hand, if we contemplate a ``timelike'' scenario, the bi-de-Sitter locked configuration might naturally explain primordial inflation. In this case, one still needs an exit mechanism (which could be provided by some instability linked to $\\gamma$ in the case where the potential $V(\\gamma)$ cannot indefinitely succeed in confining $\\gamma$).       \\vskip1cm  \\textbf{Acknowledgments:} I.K. would like to thank G.G. Ross for helpful discussions.  A.P. would like to thank R. Madden, H.P. Nilles, M. Peloso and  L. Pilo  for helpful discussions. I.K. is supported in part by PPARC rolling grant PPA/G/O/1998/00567 and  EC TMR grants HPRN-CT-2000-00152 and  HRRN-CT-2000-00148. A.P. acknowledges IHES for financial support under the Hodge Fellowship scheme.  \\vskip1cm   \\newpage          \\def\\theequation{A.\\arabic{equation}} \\setcounter{equation}{0} \\vskip0.8cm \\noindent {\\Large \\bf Appendix A: Equations of motion and constraints} \\vskip0.4cm \\noindent     The Hamiltonian of the system is: \\ba {\\mathcal H}=e^{\\sigma/2}\\left[{3 \\over 4}\\cosh\\left({\\gamma-\\delta \\over 2}\\right)(-p_{\\sigma}^2-p_{\\delta}^2+p_{\\Sigma_+}^2+p_{\\Sigma_-}^2+p_{\\Delta_+}^2+p_{\\Delta_-}^2)e^{-\\sigma}\\right.\\nonumber\\\\ \\left.+{3 \\over 2}\\sinh\\left({\\gamma-\\delta \\over 2}\\right)(-p_{\\sigma}p_{\\delta}+p_{\\Sigma_+}p_{\\Delta_+}+p_{\\Sigma_-}p_{\\Delta_-})e^{-\\sigma}+V\\right] \\ea where the canonical momenta are defined as: \\ba p_{\\Delta_\\pm}={2 \\over 3}e^{\\sigma/2}\\left[\\cosh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\Delta}_{\\pm}-\\sinh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\Sigma}_{\\pm}\\right]\\\\ p_{\\delta}={2 \\over 3}e^{\\sigma/2}\\left[\\cosh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\delta}-\\sinh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}\\right]\\\\ p_{\\sigma}={2 \\over 3}e^{\\sigma/2}\\left[\\cosh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}-\\sinh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\delta}\\right] \\ea    The equations of motion in the ``averaged proper time'' $t$ for the full Routhian are:  \\begin{itemize}   \\item $\\delta$ equation   \\be {2 \\over 3}e^{-\\sigma/2}{d \\over dt}\\left[e^{\\sigma/2}\\left(\\cosh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\delta}-\\sinh \\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}\\right)\\right]={\\de V \\over \\de \\gamma}+{\\de V \\over \\de \\delta} \\ee  \\item $\\sigma$ equation  \\be {2 \\over 3}e^{-\\sigma/2}{d \\over dt}\\left[e^{\\sigma/2}\\left(\\cosh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}-\\sinh \\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\delta}\\right)\\right]=V \\ee  \\item $\\Delta_{\\pm}$ equations   \\be e^{-\\sigma/2}{d \\over dt}\\left[p_{\\Sigma_{\\pm}}\\tanh\\left({\\gamma-\\delta \\over 2}\\right)-{2e^{\\sigma/2} \\over 3\\cosh\\left({\\gamma-\\delta \\over 2}\\right)}\\dot{\\Delta}_{\\pm}\\right]={\\de V \\over \\de \\Delta_{\\pm}} \\ee   \\item $\\gamma$ constraint   \\ba -{3 (p_{\\Sigma_+}^2+p_{\\Sigma_-}^2) \\over 8}e^{-\\sigma}{\\tanh \\left({\\gamma-\\delta \\over 2}\\right) \\over \\cosh \\left({\\gamma-\\delta \\over 2}\\right) }+{e^{-\\sigma/2} \\over 2 \\cosh^2 \\left({\\gamma-\\delta \\over 2}\\right)}(p_{\\Sigma_+}\\dot{\\Delta}_++p_{\\Sigma_-}\\dot{\\Delta}_-)\\nonumber\\\\+{1 \\over 6}\\sinh \\left({\\gamma-\\delta \\over 2}\\right) \\left[{\\dot{\\Delta}_+^2+\\dot{\\Delta}_-^2 \\over \\cosh^2\\left({\\gamma-\\delta \\over 2}\\right)}+\\dot{\\sigma}^2+\\dot{\\delta}^2\\right]-{1 \\over 3}\\cosh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}\\dot{\\delta}+{\\de V \\over \\de \\gamma}=0 \\ea  \\end{itemize}  In the above we have used the $\\gamma$ constraint to simplify the $\\delta$ equation of motion. We additionally have the Hamiltonian constraint ${\\mathcal H}=0$:   \\ba {3 (p_{\\Sigma_+}^2+p_{\\Sigma_-}^2)e^{-\\sigma}\\over 4\\cosh \\left({\\gamma-\\delta \\over 2}\\right)}+{1 \\over 3}\\cosh \\left({\\gamma-\\delta \\over 2}\\right) \\left[{\\dot{\\Delta}_+^2+\\dot{\\Delta}_-^2 \\over \\cosh^2\\left({\\gamma-\\delta \\over 2}\\right)}-\\dot{\\sigma}^2-\\dot{\\delta}^2\\right]\\nonumber\\\\+{2 \\over 3}\\sinh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}\\dot{\\delta}+V=0 \\ea     The equations of motion in the ``averaged proper time'' $t$ for the case where $p_{\\Sigma_+}=p_{\\Sigma_-}=0$ are:  \\begin{itemize}   \\item $\\delta$ equation  \\be {2 \\over 3}e^{-\\sigma/2}{d \\over dt}\\left[e^{\\sigma/2}\\left(\\cosh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\delta}-\\sinh \\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}\\right)\\right]={\\de V \\over \\de \\gamma}+{\\de V \\over \\de \\delta} \\ee  \\item $\\sigma$ equation  \\be {2 \\over 3}e^{-\\sigma/2}{d \\over dt}\\left[e^{\\sigma/2}\\left(\\cosh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}-\\sinh \\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\delta}\\right)\\right]=V \\ee  \\item $r$ equation  \\be -{2 \\over 3}e^{-\\sigma/2}{d \\over dt}\\left({e^{\\sigma/2}\\dot{r} \\over \\cosh\\left({\\gamma-\\delta \\over 2}\\right)}\\right)=-{3 p_{\\theta}^2\\over 2 r^3}e^{-\\sigma}\\cosh \\left({\\gamma-\\delta \\over 2}\\right)+{\\de V \\over \\de r} \\ee   \\item $\\gamma$ constraint   \\ba {3 p_{\\theta}^2\\over 8 r^2}e^{-\\sigma}\\sinh \\left({\\gamma-\\delta \\over 2}\\right)+{1 \\over 6}\\sinh \\left({\\gamma-\\delta \\over 2}\\right) \\left[{\\dot{r}^2 \\over \\cosh^2\\left({\\gamma-\\delta \\over 2}\\right)}+\\dot{\\sigma}^2+\\dot{\\delta}^2\\right]\\nonumber\\\\-{1 \\over 3}\\cosh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}\\dot{\\delta}+{\\de V \\over \\de \\gamma}=0 \\ea   \\end{itemize}  In the above we have used the $\\gamma$ constraint to simplify the $\\delta$ equation of motion. We additionally have the Hamiltonian constraint ${\\mathcal H}=0$:   \\ba {3 p_{\\theta}^2\\over 4 r^2}e^{-\\sigma}\\cosh \\left({\\gamma-\\delta \\over 2}\\right)+{1 \\over 3}\\cosh \\left({\\gamma-\\delta \\over 2}\\right) \\left[{\\dot{r}^2 \\over \\cosh^2\\left({\\gamma-\\delta \\over 2}\\right)}-\\dot{\\sigma}^2-\\dot{\\delta}^2\\right]\\nonumber\\\\+{2 \\over 3}\\sinh\\left({\\gamma-\\delta \\over 2}\\right)\\dot{\\sigma}\\dot{\\delta}+V=0 \\ea   Finally, the equations of motion in the ``averaged proper time'' $t$ for the case of $r=0$, when we have isotropic metrics, are:  \\ba {4 \\over 3}~{d \\over dt}\\left[e^{\\alpha+{\\gamma \\over 2}}\\dot{\\alpha}\\right]&=&e^{\\alpha+\\beta \\over 2}\\left(V-{\\de V \\over \\de \\gamma}-{\\de V \\over \\de \\delta}\\right)\\\\ {4 \\over 3}~{d \\over dt}\\left[e^{\\beta-{\\gamma \\over 2}}\\dot{\\beta}\\right]&=&e^{\\alpha+\\beta \\over 2}\\left(V+{\\de V \\over \\de \\gamma}+{\\de V \\over \\de \\delta}\\right) \\ea  We additionally have the two Friedmann constraints (which we have also used to simplify the above equations of motion):  \\ba {2 \\over 3}\\dot{\\alpha}^2&=&e^{-{\\gamma-\\delta \\over 2}}\\left({V \\over 2}-{\\de V \\over \\de \\gamma}\\right)\\\\ {2 \\over 3}\\dot{\\beta}^2&=&e^{{\\gamma-\\delta \\over 2}}\\left({V \\over 2}+{\\de V \\over \\de \\gamma}\\right) \\ea     \\newpage       \\def\\theequation{B.\\arabic{equation}} \\setcounter{equation}{0} \\vskip0.8cm \\noindent {\\Large \\bf Appendix B: Perturbation analysis around the $b^{\\mu}=0$ shift vector in the cosmological metric ansatz} \\vskip0.4cm \\noindent    In this Appendix we study under what conditions there might exist more solutions than the ``trivial'' solution $b^{\\mu}=0$ discussed in the text. The eigenvalue problem: \\be (g_1^{-1}g_2)^{M}_{N}e^{N}_A=\\lambda_A e^{M}_A \\ee considered for $b^{\\mu} \\neq 0$ but small,  can be related (to ${\\cal O}(b^2)$) by standard techniques to the unperturbed  one (denoted by an overbar)  as: \\ba \\lambda_0&=&\\bar{\\lambda}_0\\left(1-\\sum_a {\\bar{\\lambda}_a (b^a)^2 \\over \\bar{\\lambda}_0-\\bar{\\lambda}_a}\\right)\\\\ \\lambda_a&=&\\bar{\\lambda}_a\\left(1+ {\\bar{\\lambda}_a (b^a)^2 \\over \\bar{\\lambda}_0-\\bar{\\lambda}_a}\\right) ~~~ ({\\rm no~sum}) \\ea where $b^a=\\bar{e}^a_{\\mu}b^{\\mu}$. The variation of the action when making a  perturbation is: \\ba \\delta S_m&=&-{1 \\over 2}\\sqrt{-g_2}T^{(2)~A}_{~~~A}~{\\delta\\lambda_A \\over \\lambda_A}\\\\ &=&\\sqrt{-g_2}\\sum_a  {\\bar{\\lambda}_a b^a \\over \\bar{\\lambda}_0-\\bar{\\lambda}_a}(T^{(2)~0}_{~~~0}-T^{(2)~a}_{~~~a})\\delta b^a \\ea and hence: \\be {\\delta S_m \\over \\delta b^a}=\\sqrt{-g_2} {\\bar{\\lambda}_a b^a \\over \\bar{\\lambda}_0-\\bar{\\lambda}_a}(T^{(2)~0}_{~~~0}-T^{(2)~a}_{~~~a})~~~ ({\\rm no~sum})\\label{bvar} \\ee  As said in the text, the equation of motion of $b^{a}$ is the vanishing of (\\ref{bvar}). There are only two possible ways to make the above quantity vanish: either $b^a=0$ or $T^{(2)~0}_{~~~0}-T^{(2)~a}_{~~~a}=0$. The first possibility means that $b^a=0$ is an isolated solution. Therefore, it is the second possibility which signals the threshold for the existence of new solutions, besides the trivial one. [We are assuming here that, as some parameters vary, the nonperturbative solutions can be made to coincide with the perturbative one.] On the other hand according to (\\ref{energy1}),(\\ref{energy2}): \\be T^{(2)~0}_{~~~0}-T^{(2)~a}_{~~~a}=-2e^{-\\sigma_1/4}(\\de_{\\mu_0}V-\\de_{\\mu_a}V)=-4e^{-\\sigma_1/4}(\\mu_0-\\mu_a)\\de_{\\sigma_2}V \\ee for potentials of the class $V=V(\\sigma_1,\\sigma_2)$. Therefore, a general necessary condition for the possible existence of non-perturbative solutions is that $\\de_{\\sigma_2}V = 0$ admits solutions.       \\def\\theequation{C.\\arabic{equation}} \\setcounter{equation}{0} \\vskip0.8cm \\noindent {\\Large \\bf Appendix C: Analytic solutions for the Pauli-Fierz potential in the timelike worldline separation limit} \\vskip0.4cm \\noindent    In this Appendix we sill study an analytic description of the solutions of the extreme timelike worldline separation  ({\\textit{i.e.}} $r=0$) for the original Pauli-Fierz potential (\\ref{PF}). As discussed in the text, the solutions exhibit an initial stage of acceleration for large $\\delta$ and a period of deflation as $\\delta \\to 0$ for the first metric, while the second metric remains approximately flat.  We will split up the analysis of this system into the two above-mentioned asymptotic regions.  Let us note  the equations of state for the two metrics obtained by (\\ref{state1}), (\\ref{state2}): \\ba w_1={P_1 \\over \\rho_1}=-1-2~{\\delta-3\\gamma \\over \\delta^2-6\\delta+3\\delta \\gamma}\\label{isostate1}\\\\ w_2={P_2 \\over \\rho_2}=-1+2~{\\delta-3\\gamma \\over \\delta^2+6\\delta+3\\delta \\gamma}\\label{isostate2} \\ea   \\begin{itemize}  \\item {\\bf The cosmological initial singularity limit where   $\\gamma \\to -\\infty$ and $\\delta \\to \\infty$}  \\end{itemize}  We have already qualitatively described in the text the behaviour of the system in this limit. Due to this motion in field space, we can see numerically that for all initial conditions, the first metric experiences accelerated expansion. As an example to illustrate the behaviour of the system in this limit, we will consider the case where $\\gamma$ initially lies very near the $\\gamma=-2-{1 \\over 3}\\delta$ line. Then, the solution follows this line in a very good approximation both back in time towards the initial singularity, as well as forward in time until $\\delta \\sim {\\mathcal{O}}(1)$. From the $\\beta$ Friedman equation  we get that $\\dot{\\beta}\\approx 0$ and from the $\\alpha$ constraint that: \\be \\dot{\\alpha}^2=-3m^2(\\gamma+2)e^{-(2\\gamma+3)} \\ee But approximately $\\dot{\\alpha}=-\\dot{\\delta}=3\\dot{\\gamma}$, so we have a differential equation for $\\gamma$. Integrating this we get: \\ba {\\rm Erf}(\\sqrt{-(\\gamma+2)})=1-{t\\over t_{cr}}&\\\\ \\Rightarrow~~\\gamma=-2-\\left[{\\rm Erf^{(-1)}}\\left(1-{t\\over t_{cr}}\\right)\\right]^2 &~~\\to~~\\log\\left(t \\over t_{cr}\\right) \\ea where $t_{cr}={1\\over m}\\sqrt{3\\pi \\over e}$, ${\\rm Erf}(x)={2 \\over \\sqrt{\\pi}}\\int_0^x e^{-y^2}dy$ is the error function, ${\\rm Erf^{(-1)}}(x)$ the inverse error function. The limit we have indicated is at $t \\to 0$.  From the properties of the error function we have that if $t \\to 0^+$ then $\\gamma \\to -\\infty$. In the other limit that  $t \\to t_{cr}$ we get that $\\gamma \\to -2$. However, we never reach the latter limit, since close to that, our approximation breaks down because $\\delta \\sim {\\mathcal{O}}(1)$. In the region that this approximation is valid, we have: \\be \\delta=-3(\\gamma+2)~~,~~\\sigma=-\\delta+2C~~,~~\\alpha=-\\delta+C~~,~~\\beta=C \\ee where $C$ is an integration constant. We see that the first metric expands, but one should describe this expansion in this metric's proper time. Asymptotically, for $t \\to 0$ we have: \\be dt_1=e^{-\\gamma/2}dt~~\\Rightarrow~~t_1=2\\sqrt{t_{cr}t} \\ee On the other hand the proper time of the second metric is: \\be dt_2=e^{\\gamma/2}dt~~\\Rightarrow~~t_2={2 \\over 3 \\sqrt{t_{cr}}}t^{3/2} \\ee Thus, the first metric is  intrinsically inflating with the scale factor behaving as  $e^{\\alpha/3} \\sim t_1^2$, while the second metric is approximately flat.  The equation of state (\\ref{isostate1}), in  this limiting case which we are examining, is for the first metric: \\be w_1=-1+{1 \\over 3}~{\\gamma+1 \\over \\gamma+2}~~\\to~~-{2 \\over 3} \\ee while for the second one, (\\ref{isostate2}) leads to  $w_2 \\to -\\infty$, since the  $\\gamma=-2-{1 \\over 3}\\delta$ line is the root of the denominator.  Let us now check if the obtained solution is stable against perturbations of $r$. The variation of the Routhian in quadratic order is: \\be \\delta {\\mathcal L}={1 \\over 3}e^{\\sigma /2}\\left({\\dot{r}^2 \\over \\cosh\\left(\\gamma-\\delta \\over 2\\right)}-{1 \\over 3}r^2\\right) \\ee From the extremization of this action we get the following motion for $r$: \\be r=C_1J_0\\left({1 \\over 6}\\sqrt{e \\over 2\\pi} m^2t^2\\right)+C_2Y_0\\left({1 \\over 6}\\sqrt{e \\over 2\\pi} m^2 t^2\\right) \\ee which is growing only logarithmically as $t \\to 0$. Thus, the solution is stable in very good approximation. This is in agreement with our numerical study.    \\begin{itemize}  \\item {\\bf The final state limit where $\\gamma \\to -\\infty$, $\\delta \\to 0$}  \\end{itemize}  The final state of the evolution of the system is independent of the initial conditions. As $\\delta \\to 0$ we find that  $\\gamma \\to -\\infty$ and the first metric deflates and becomes asymptotically flat in infinite proper time, while the second experiences a finite (in proper time) period of inflation. The runaway of $\\gamma$ in this limit is an unavoidable fact of the Pauli-Fierz potential, as we have discussed in the text, because the action ceases to have an extremum at finite $\\gamma$, when one of the two worldline velocities tends to zero. In order to study this limit we need to do a different approximation to the equations of motion. Combining the two Friedman constraints (\\ref{F1}), (\\ref{F2}) and keeping leading terms we get the following relation: \\be {\\dot{\\alpha}^2 \\over \\dot{\\beta}^2}=e^{-\\gamma} \\label{ratio} \\ee The equations of motion on the other hand are approximated by: \\ba {4 \\over 3}{d \\over dt}\\left(e^{\\alpha+\\gamma/2}\\dot{\\alpha}\\right)={\\gamma \\over 3}e^{\\alpha+\\beta \\over 2}\\\\ {4 \\over 3}{d \\over dt}\\left(e^{\\beta-\\gamma/2}\\dot{\\beta}\\right)=-{\\gamma \\over 3}e^{\\alpha+\\beta \\over 2} \\ea Using (\\ref{ratio}) we have: \\ba {d \\over dt}\\left(e^{\\alpha}\\dot{\\beta}\\right)={m^2 \\over 2}e^{\\alpha+\\beta \\over 2}\\log{\\dot{\\beta} \\over \\dot{\\alpha}}\\\\ {d \\over dt}\\left(e^{\\beta}\\dot{\\alpha}\\right)=-{m^2 \\over 2}e^{\\alpha+\\beta \\over 2}\\log{\\dot{\\beta} \\over \\dot{\\alpha}} \\ea Then since $\\dot{\\alpha}\\dot{\\beta} \\to 0$, we can neglect this term and obtain the system: \\ba \\ddot{\\beta}={m^2 \\over 2}e^{\\beta-\\alpha\\over 2}\\log{\\dot{\\beta} \\over \\dot{\\alpha}} \\approx {m^2 \\over 2}\\log{\\dot{\\beta} \\over \\dot{\\alpha}}\\\\ \\ddot{\\alpha}=-{m^2 \\over 2}e^{-{\\beta-\\alpha \\over 2}}\\log{\\dot{\\beta} \\over \\dot{\\alpha}}\\approx -{m^2 \\over 2}\\log{\\dot{\\beta} \\over \\dot{\\alpha}} \\ea From this system we get: \\be \\dot{\\beta}=\\dot{\\alpha}\\left({C_1 \\over \\dot{\\alpha}}-1\\right) \\ee which putting back into the second equation gives: \\be \\ddot{\\alpha}=-{m^2 \\over 2}\\log\\left({C_1 \\over \\dot{\\alpha}}-1\\right) \\ee This can be solved in the region where $\\dot{\\alpha} \\to C_1$ and gives: \\be \\dot{\\alpha}={C_1 \\over 1+{\\rm li}^{(-1)}\\left({m^2  \\over 2C_1}(t-t_0)\\right)} \\approx C_1 \\left[1-{m^2  \\over 2C_1}(t-t_0)\\log \\left({m^2  \\over 2C_1}|t-t_0|\\right)\\right] \\ee where ${\\rm li}(x)=\\int_0^x {dy \\over \\log y}$ is the logarithmic integral and ${\\rm li}^{(-1)}(x)$ its inverse function. The limits which have been used are that for $x \\to 0^+$, ${\\rm li}(x) \\to {x \\over \\log x}$ and for $x \\to 0^-$, ${\\rm li}^{(-1)}(x) \\to x \\log |x|$. This shows that $\\ddot{\\alpha}$ diverges to $-\\infty$. Asymptotically, the function  $\\alpha$ is: \\be \\alpha \\approx C_2+C_1 (t-t_0)-{m^2 \\over 4}(t-t_0)^2\\log \\left({m^2  \\over 2C_1}|t-t_0|\\right) \\ee The asymptotics for the $\\beta$ function is: \\be \\beta=C_2+{m^2 \\over 4}(t-t_0)^2\\log \\left({m^2  \\over 2C_1}|t-t_0|\\right) \\ee and for $\\delta$: \\be \\delta=-C_1 (t-t_0) \\ee On the other hand $\\gamma$ is: \\ba \\gamma=2\\log \\left({C_1 \\over \\dot{\\alpha}}-1\\right)~~\\Rightarrow~~\\gamma&=&2\\log \\left[{\\rm li}^{(-1)}\\left({m^2  \\over 2C_1}(t-t_0)\\right)\\right]\\nonumber\\\\&\\approx& 2\\log \\left[{m^2  \\over 2C_1}(t-t_0)\\log \\left({m^2  \\over 2C_1}|t-t_0|\\right)\\right] \\ea   The proper time in the first metric is: \\be dt_1=e^{-\\gamma/2}dt~~\\Rightarrow~~t_1 \\sim {2C_1 \\over m^2}\\log \\left|\\log\\left({m^2 \\over 2C_1}|t-t_0|\\right)\\right| \\ee thus, the singularity point $t_0$ in the proper time of the first metric is at $t_1 \\to \\infty$. On the other hand, the proper time for the second metric is: \\be dt_2=e^{-\\gamma/2}dt~~\\Rightarrow~~t_2 = t_{20}  + {2C_1 \\over m^2}\\left[{m^2 \\over 2C_1}(t-t_0)\\right]^2 \\log \\left[{m^2 \\over 2C_1}|t-t_0|\\right] \\ee and it is finite at $t \\to t_0$.  Thus, the first metric is intrinsically deflating  with the scale factor behaving as: \\be e^{\\alpha/3} \\sim \\left(1-{2C_1^2 \\over 3 m^2}e^{-e^{m^2t_1 \\over 2C_1}}\\right) \\ee while the second is  intrinsically inflating as: \\be e^{\\beta/3} \\sim e^{{C_1 \\over 6}t_2} \\ee  The equation of state (\\ref{isostate1}) on the first metric, in this limit, is: \\be w_1=2~ {1 \\over \\delta} ~~\\to~~+\\infty \\ee while on the second one (\\ref{isostate2}) leads to the opposite effect: \\be w_2=-2~ {1 \\over \\delta} ~~\\to~~-\\infty \\ee       \\def\\theequation{D.\\arabic{equation}} \\setcounter{equation}{0} \\vskip0.8cm \\noindent {\\Large \\bf Appendix D: The general $p \\neq 0$ case evolution for the brane motivated potential} \\vskip0.4cm \\noindent    In this Appendix we will discuss the evolution of the system for the brane motivated potential for the general case where the ``initial kinetic energy'' $p$ is non-zero. The solution of the equations of motion is given by (\\ref{x2sol}), (\\ref{x1sol}). The initial ``incoming'' solutions are then for $\\tau \\ll 0$ where the kinetic energy of the system is large in comparison with the potential energy. Then  we can write the solutions of $\\sigma$ and $r$ using (\\ref{x1x2rs}) and separating the various constants (and subleading terms) as: \\begin{itemize} \\item For $p>0$ \\be \\sigma \\to {3 \\over 7}(3-\\sqrt{2})p\\tau + \\cdots ~~~,~~~r \\to \\sigma +\\cdots \\ee \\item For $p<0$ \\be \\sigma \\to -{3 \\over 7}(3+\\sqrt{2})p\\tau + \\cdots ~~~,~~~r \\to -\\sigma +\\cdots \\ee \\end{itemize} On the other hand, the final state solutions are  for $\\tau \\to 0$ when, due to the expansion of the two metrics, the kinetic energy of the system has become subdominant in comparison with the potential energy. Then  we have asymptotically for both signs of $p$: \\be \\sigma \\to -{18 \\over 7}\\log |p\\tau| + \\cdots ~~~,~~~r \\to -{\\sqrt{2} \\over 3}\\sigma +\\cdots \\ee Note that the solutions in the latter epoch,  have the same scaling law as the ones for $p=0$. In the above language we see that in the case where $p>0$, $r$ initially increases until a maximum value and then shrinks to zero. On the other hand for $p < 0$ $r$ always increases. In all cases $\\sigma$ increases and thus the volume of each metric expands.  Now, we need  to go back to proper time $t$ to see the behaviour of our solutions. For $\\tau \\ll 0$ both cases have $\\sigma = c\\tau+\\cdots$ with different constants $c$, whose value does not have any significance as we will see in the following. Then the proper time and the expression of $\\sigma$ as a function of it are: \\be t = {2 \\over c}~e^{c\\tau+\\cdots \\over 2}~~~,~~~ \\sigma=2 \\log (p t) +\\cdots \\label{pastasym} \\ee where again we ignored unimportant constants. By rescaling coordinates to absorb the latter constants we finally see that each metric has a {\\textit{power-law behaviour}}, and if we parametrise it in the standard way as: \\be ds^2=-dt^2+\\sum_{\\mu=1}^3(p t)^{2 p^{\\mu}}(dx^{\\mu})^2 \\label{Kas1} \\ee we have that the Kasner exponents $p^{\\mu}$ for $p>0$ are: \\be \\renewcommand{\\arraystretch}{1.5} p^{\\mu}_{(1)}=\\left[\\begin{array}{ccc}{1 \\over 3}+{1 \\over \\sqrt{3}}\\\\{1 \\over 3}-{1 \\over \\sqrt{3}}\\\\{1 \\over 3}\\end{array}\\right]~~~{\\rm and}~~~p^{\\mu}_{(2)}=\\left[\\begin{array}{ccc}{1 \\over 3}-{1 \\over \\sqrt{3}}\\\\{1 \\over 3}+{1 \\over \\sqrt{3}}\\\\{1 \\over 3}\\end{array}\\right] \\ee and for $p<0$ they are the same with a flip on the sign of the second addendum of the first two exponents.  The above exponents satisfy the usual quadratic Kasner relation $\\sum_{\\mu} (p^{\\mu})^2-\\left(\\sum_{\\mu} p^{\\mu}\\right)^2=0$, as well as  $\\sum_{\\mu} p^{\\mu}=1$, which means that each metric's volume $v_i$ expands as a function of its respective proper time as $v \\propto t_i$, with $i=1,2$.  Thus, this evolution is highly anisotropic with: \\be A_1=A_2=\\sqrt{2} \\ee  One can notice at this point a potential paradox because an exact Kasner metric is known to be ``on the light cone'', {\\textit{i.e.}} to  have $\\dot{\\alpha}^{\\mu}\\dot{\\alpha}_{\\mu}={1 \\over 6}(\\dot{\\sigma}^2-\\dot{r}^2)=0$. On the other hand, from the Hamiltonian  constraint (\\ref{eq_H}) the same quantity should be very large for $r \\gg 1$. This can be understood if we include next to leading order terms in our asymptotic solution. These will modify the $\\sigma$ asymptotic (\\ref{pastasym}) by a term linear in $t$ and also the $r$ asymptotic with a term linear in $t$ with different coefficient. This will immediately render  $\\dot{\\sigma}^2-\\dot{r}^2$ very large as expected.  For the case where $\\tau \\to 0$, we can express the proper time and $\\sigma$ as a function of it as: \\be t=C \\tau^{-{2 \\over 7}}~~~,~~~ \\sigma=9 \\log (p t) +\\cdots \\ee Thus, the exponents this time are: \\be \\renewcommand{\\arraystretch}{1.5} p^{\\mu}_{(1)}=\\left[\\begin{array}{ccc}{3 \\over 2}-\\sqrt{{3 \\over 2}}\\\\{3 \\over 2}+\\sqrt{{3 \\over 2}}\\\\{3 \\over 2}\\end{array}\\right]~~~{\\rm and}~~~p^{\\mu}_{(2)}=\\left[\\begin{array}{ccc}{3 \\over 2}+\\sqrt{{3 \\over 2}}\\\\{3 \\over 2}-\\sqrt{{3 \\over 2}}\\\\{3 \\over 2}\\end{array}\\right] \\ee The above exponents do not satisfy the quadratic (zero-mass-shell) Kasner relation, since  $\\sum_{\\mu} (p^{\\mu})^2-\\left(\\sum_{\\mu} p^{\\mu}\\right)^2=-{21 \\over 2}$. In addition, $\\sum_{\\mu} p^{\\mu}={9 \\over 2}$, which means that each metric's volume $v_i$ expands as a function of its respective proper time as $v \\propto t_1^{9/2}$, with $i=1,2$. Note that the latter volume expansion has an accelerating behaviour and is exactly the same as in the $p=0$ case.   The evolution is still anisotropic but slightly less than the Kasner case: \\be A_1=A_2={2 \\over 3} \\ee  Let us now check, as we did for the $p=0$ case,  if the obtained solution is stable against perturbations of $\\delta$. We have to recalculate the quantities $A$ and $B$ of Eq.(\\ref{ABdef})  and their limits for  $r \\gg 1$. For $\\tau \\ll 0$ both signs of $p$ have: \\be A=0 ~~~,~~~B\\to {m^4 t^2 \\over 12}e^{\\sigma/2} \\ee and hence we have $\\delta=0$ and absolutely stable motion. On the other hand, for the case $\\tau \\to 0$, since the evolution is exactly the same as for the $p=0$ case, the solution is unstable at late times as described in the main text. This instability is again linked to a run away of $\\gamma$ towards large values."
    },
    "0206/astro-ph0206473_arXiv.txt": {
        "abstract": "We present high-resolution spectroscopic observations of the \\ion{Li}{1} resonance line in a sample of 62 stars that belong to 31 common-proper motion pairs with twin F or G-type components. Photospheric abundances of lithium were derived by spectral synthesis analysis. For seven of the pairs, we have measured large lithium abundance differences. Eleven other pairs have components with similar lithium abundances.  We cannot determine if the remaining 13 pairs have lithium differences because we did not detect the \\ion{Li}{1} lines, and hence we can only provide upper limits to the abundances of both stars. Our results demonstrate that twin stars do not always share the same lithium abundances. Lithium depletion in solar-type stars does not only depend on age, mass, and metallicity. This result is consistent with the spread in lithium abundances among solar-type stars in the solar-age open cluster M67. Our stars are brighter than the M67 members of similar spectral type, making them good targets for detailed follow-up studies that could shed light on the elusive mechanism responsible for the depletion of lithium during the main-sequence evolution of the Sun and solar-type stars. ",
        "introduction": "Lithium has a rich life. It is created in many environments (primordial nucleosynthesis, cosmic ray spallation in the interstellar medium, spallation and fusion reactions around compact relativistic objects, flares, and red giant thermal pulses), and it is destroyed in the interior of stars by collision with protons at a temperature of about 2.7$\\times 10^6$~K. The Sun has a lithium abundance more than two orders of magnitude lower than meteoritic material (M\\\"uller, Peytremann, \\& de la Reza 1975), which reflects the composition of the presolar nebula. The Sun probably did not deplete lithium during the pre main-sequence evolution, but rather during the slow main-sequence evolution, although the exact mechanism has not been identified yet (Mart\\'\\i n 1997, 1998). Recent models of lithium depletion in solar-type stars suggest that two types of mixing may be at work during main-sequence evolution; namely overshooting and rotational mixing (Umezu \\& Saio 2000). The history of rotational mixing may be strongly influenced by the initial conditions of angular momentum distribution in the protoplanetary disk. Thus, it is conceivable that there could be a connection between the rate of lithium depletion by rotational mixing and the presence of companions (stellar or substellar) to the stars.   Binary stars with twin components are interesting for understanding binary formation and evolution. Common proper motion (CPM) pairs of twins allow us a test of the validity of stellar evolution models because their separations are so large ($>$100 AU) that each member has probably evolved independently. If all the properties of stars are determined by age, mass, and chemical composition, twins should be identical. If the lithium abundance of a main-sequence star is determined by its age, mass, and chemical composition, pairs of twins should show the same lithium abundances. In the early work of Herbig (1965), it was already apparent that there was a scatter in lithium abundances among pairs of twins. Herbig's study included 53~UMa~A and B, where lithium was detected in A but not in B; 16~Cyg A and B, where lithium was detected in A but not in B; and 53 Aqr A and B, where lithium was detected in both stars.   The system 16~Cyg has received some attention recently. The two stars of the pair are nearly identical. The primary has V=5.96, T$_{\\rm eff}$=5785~K and log~g=4.28. The secondary has V=6.20, T$_{\\rm eff}$=5747~K and log~g=4.35. Despite those similarities, the members of this pair have two fundamental differences. King et al. (1997) have shown that the primary has a lithium abundance higher than that of the Sun (logN(Li)=1.27$\\pm$0.05 in the customary scale of log~N(H)=12), while the secondary has a much lower lithium abundance (logN(Li)$<$0.60). The large difference in lithium abundance is surprising because both stars should be the same age and have very similar masses. Another significant difference between the two stars is that the secondary harbors a giant planet  with a minimum mass of 1.5 Jupiters, an eccentricity of e=0.63 and a period of 800.8 days (Cochran et al. 1997). On the other hand, no giant planet has yet been detected around 16~Cyg~A. However, such nondetection does not rule out the presence of giant planets with long periods or smaller planets with short periods.  Gonzalez \\& Laws (2000) and Ryan (2000) have studied the distribution of lithium abundaces among planet-harboring stars, and have reached opposite conclusions. While Gonzalez \\& Laws claim that lithium is overdepleted in stars with giant planets, Ryan claims that it is normal. The disagreement comes from the choice of stars to be compared with the host stars of extrasolar planets. This is a difficult task because the ages of most field stars are not well known.  Pairs of twins offer the possibility of comparing stars that are coeval, and have nearly the same mass and metallicity. Thus, we may expect that the systematic study of binaries may shed light on whether there is any connection between the presence of giant planets and lithium depletion. Gratton et al. (2001) have recently reported an abundance analysis for six wide binaries. They find that the primary star of the binary HD~219542 has higher iron and lithium abundances than the secondary star, which leads them to suggest that the primary has ingested a planet.  In this paper we report lithium abundances for a sample of 31 visual binaries. The targets have been chosen to be as similar as possible to 16~Cyg in spectral type and magnitude. The paper is organized as follows: In Section 2 we describe the observations. In Section~3 we present the abundance analysis. In Section~4 we discuss each pair of twins individually when we could detect the \\ion{Li}{1} line in at least one of the stars, and we compare them with lithium abundances in open clusters and theoretical models. Section~5 contains the discussion, and in Section~6 we present our conclusions. ",
        "conclusions": "Lithium abundance is very sensitive to the mass (and thus T$_{\\rm eff}$) of stars. In the temperature range of our program stars (6500-5250~K), lithium abundances generally decrease with decreasing T$_{\\rm eff}$, although at the hot end of this temperature range, the F-type lithium dip starts to kick in (Boesgaard \\& Tripicco 1987). If the lithium abundances of the components of pairs of twins were dominated by the mass dependance of lithium depletion, we would expect a correlation between the difference in T$_{\\rm eff}$ and the difference in logN(Li), in the sense that the cooler components should show systematically lower lithium abundances than the hotter ones. Figure~6 shows that such correlation is not present in the data. We conclude that the observed pattern of lithium abundances in pairs of twins cannot be explained solely with the dependance of lithium depletion on stellar mass.  We have found five pairs with nearly identical T$_{\\rm eff}$ and nearly identical logN(Li), and four pairs with nearly identical T$_{\\rm eff}$ and very different logN(Li). Hereafter we call these latter four pairs \"16~Cyg~analogs\", namely: HD~6872 A/B, BD+60~0269/BD+60~0271, BD+13~2311~A/B, and HD 98744/HD 98745. The common characteristics of \"16~Cyg~analogs\" is that they are pairs of twins with lithium abundance difference between the two stars of the pair that exceeds 0.5 dex (which corresponds to 3~$\\sigma$ significance in our synthetic fit analysis) and difference in T$_{\\rm eff}<$200~K.  Figure~7 illustrates the comparison between the pairs of twins and the Hyades cluster in a T$_{\\rm eff}$ versus logN(Li) diagram. The Hyades single star members define a tight relationship between T$_{\\rm eff}$ and logN(Li). In general the twins do not follow the Hyades relation. A possible explanation for this discrepancy could be that there are many tidally locked binaries (TLB) hidden among the twins. We do not know of any example of a TLB in our sample.  HD~8610 has a lithium abundance similar to Hyades members with the same T$_{\\rm eff}$. The lithium-T$_{\\rm eff}$ locus in the Hyades declines steeply from 5500~K to 5250~K, so one would expect that HD~8624 should have a lithium abundance much lower than that of HD~8610, which is contrary to what we have found. HD~8624 is a double-lined spectroscopic binary with a period of 14.91 days (Tokovinin 1999) and an eccentricity e=0132. Even though it is not strictly a TLB because the orbit is not completely circularized, it may have preserved lithium in a manner similar to the TLBs studied by Barrado y Navascu\\'es et al. (1997), which show a trend of higher lithium abundances when compared to single stars. However, TLBs are relatively rare among the general population of solar-type stars. A program to monitor the radial velocity of 10 of our pairs has been started using HIRES with the iodine cell. There are already enough data for 16 of the stars to find TLBs, but none has been seen (Marcy 2002, private communication). We conclude that it is very unlikely that TLBs can account for most of the discrepancy between the pairs of twins and the Hyades.  The average age of the pairs of twins is probably larger than the Hyades cluster age ($\\sim$600~Myr). It is informative to compare the twins with the well-studied solar-age cluster M67. In Figure~7 we compare the pairs of twins with the members of M67.  The M67 stars show a large spread of lithium abundances at a given T$_{\\rm eff}$, which is similar to the behavior exhibited by our binaries. The evolution of lithium from the Hyades to M67 implies a mechanism of lithium depletion during the main-sequence lifetime of solar-type stars that is not very mass dependent. We have not found in the literature any models that can explain the evolutionary pattern of lithium abundances implied by the comparison between the Hyades and the M67 solar-type stars. We conclude that there is currently not a good theoretical explanation for the lithium abundances observed in our pairs of twins and the open clusters. In future papers we plan to tackle this problem observationally via searches for connections between the lithium abundances in the pairs of twins and the presence of companions (stellar or substellar), chromospheric activity, rotation, and the abundances of other chemical elements. At the moment it would be premature to suppose that the presence of planets had anything to do with the spread in lithium abundances."
    },
    "0206/astro-ph0206190_arXiv.txt": {
        "abstract": "{We use data from a recent long ASCA observation of the Narrow Line Seyfert 1 galaxy \\object{Ark~564} to investigate in detail its timing properties. We show that a thorough analysis of the time series, employing techniques not generally applied to AGN light curves, can provide useful information to characterize the engines of these powerful sources. We searched for signs of non--stationarity in the data, but did not find strong evidences for it. We find that the process causing the variability is very likely nonlinear, suggesting that variability models based on many active regions, as the shot noise model, may not be applicable to \\object{Ark~564}. The complex light curve can be viewed, for a limited range of time scales (as indicated by the breaks in the structure and power density spectrum), as a fractal object with non--trivial fractal dimension and statistical self--similarity. Finally, using a nonlinear statistic based on the scaling index as a tool to discriminate time series, we demonstrate that the high and low count rate states, which are indistinguishable on the basis of their autocorrelation, structure and probability density functions, are intrinsically different, with the high state characterized by higher complexity. ",
        "introduction": "Active Galactic Nuclei (AGN) are variable in every observable wave band. The X-ray flux exhibits variability on time scales shorter than any other energy band, indicating that the emission occurs in the innermost regions of the central engine. Therefore, a study of the X-ray variability provides an additional powerful tool to probe the extreme physical processes operating in the inner parts of the accretion flow close to the accreting black hole. Although X-ray variability has been observed in AGN for more than two decades, its origin and nature is still poorly understood. Until recently, there were no sufficiently long observations of AGN with good signal to noise to allow a thorough analysis. Secondly, the temporal analysis is frequently limited to calculations in the Fourier domain of the power density spectrum (or, equivalently, in the time domain of the structure function) which, although useful for detecting possible periodicities or typical time scales in a signal, does not exploit all the potential information contained in a time series. The reason is that this technique uses only the first two moments (namely the mean and the variance) of the probability distribution function associated with the physical processes underlying the signal. However, only a Gaussian distribution can be completely described by the first two moments, and there are several indications that probability density functions associated with X-ray light curves both in AGN and in Galactic black hole systems, are not Gaussian (e.g. Leighly \\cite{leig1}, Greenhough \\cite{green}). Finally, the light curves, i.e. the starting point for any kind of temporal analysis, have been often considered as a by-product of the spectral analysis, which still catalyzes most of the attention and efforts, in particular now that high resolution X-ray spectroscopy of AGN is possible thank to {\\it Chandra} and {\\it XMM-Newton}. Spectral analysis proved to be very useful in providing constraints on the physical parameters of the accretion flow around black holes but it must be pointed out that, due to low signal-to-noise data, spectral models are often applied to time-averaged spectra, despite the fact that sources show rapid variability. Thus the additional information provided by timing observations can be crucial to break the degeneracy among spectral models.  One of the most critical open questions, related to the X-ray variability in AGN, concerns the nature of the variability: is it linear or nonlinear? In a mathematical sense linearity means that the value of the time series at a given time can be written as a linear combination of the values at previous times plus some random variable. A positive detection of nonlinearity would have an immediate and very important consequence for the modeling of the region producing X-rays: all the variability models based on many independent active regions, as the shot noise model (Terrel 1972) or magnetic flares (e.g. Galeev et al. 1979) would be ruled out in favor of inherently nonlinear models as the self-organized criticality disk model (Mineshige et al. 1994) or the emission of X--ray radiation from a putative jet.  \\begin{figure} \\psfig{figure=f1.ps,height=6.5cm,width=8.7cm,% bbllx=124pt,bblly=153pt,bburx=692pt,bbury=581pt,angle=0,clip=} \\caption{ASCA SIS0 light curve of \\object{Ark564}. The dot--dashed thick line divides the light curves in two halves used to investigate stationarity. The dashed lines define two intervals used to characterize the timing behavior during the low and high count rate states. \\label{figure:lc-ark}} \\end{figure} In an attempt to answer this question and, more generally, to investigate the nature of the X-ray variability, we focus on a prominent object of a particular class of AGN, the Narrow-Line Seyfert 1 galaxies (NLS1), which often display rapid, large amplitude X-ray variability as well as extreme long-term changes (Forster \\& Halpern \\cite{forst}, Boller \\cite{boll2}, Brandt \\cite{brand2}), and therefore represent the ideal objects for an X-ray temporal analysis. In particular we analyze a recent long ASCA observation of \\object{Ark~564}, the brightest NLS1 in the 2-10 keV band, using non-standard (at least for X-ray astronomy) analysis techniques.  The outline of this paper is the following. In Sect. 2 we describe the data used for the timing analysis. In Sect. 3 we investigate the important issue of the stationarity of a time series, by dividing the light curve into two equal parts and by computing and comparing the mean, the variance, the auto-correlation and structure functions and the power spectra of the two halves. Sect. 4 deals with the search for nonlinearity. To this end, we choose two parts of the light curve with reasonable length ($\\sim$ 4 days), when the source was at a high and low count rate state, respectively. We use a new technique based on the constrained randomization of a time series and the nonlinear prediction error as an indicator of nonlinearity, in order to test the hypothesis that the light  curve is a realization of a linear process. In Sect. 5 we first investigate the variability behavior of the source, utilizing both standard (excess variance, probability density function) and non--standard (fractals) techniques. Then, in order to investigate further whether the behavior of the source is identical in the two states (high and low) we used techniques like the phase space reconstruction and the scaling index method. Finally, in Sect. 6 we draw the main conclusions.  \\begin{figure} \\psfig{figure=h3233f2.ps,height=12cm,width=8.7cm,% bbllx=44pt,bblly=29pt,bburx=370pt,bbury=567pt,angle=0,clip=} \\caption{ASCA SIS0 light curves of \\object{Ark564} during the high (top panel) and the low (bottom panel) count rate state. Time bins are 16s. Both light curves contain 10000 data points each. \\label{figure:lc2a}} \\end{figure} ",
        "conclusions": "We have carried out a thorough analysis of the timing properties of the NLS1 \\object{Ark~564}, using the data from a recent long observation performed by the ASCA satellite.  The first important result is that we searched for signs of non--stationarity, without finding any strong evidence for it, although the differences in the structure function and autocorrelation functions between the first and the second part of the light curve give some indications in favor of a non--stationarity of the total light curve. To asses the stationarity of a time series is also important for a different aspect: the non-stationarity of the signal can strongly affect the identification of genuine nonlinearity possibly present in the data, and lead to wrong interpretations of results obtained from nonlinear statistical analyses.  As sensitive light curves long enough to cover several transitions between high and low states are presently not available, we concentrated on two sub--intervals of the \\object{Ark~564} light curve, which are locally stationary, contain a sufficiently large number of data points (10000, using a bin size of 16 s) for a meaningful statistical analysis, and which have significantly different mean count rates in order to characterize the intrinsic temporal differences from the low and the high state of the source.  We then carefully addressed the issue of the possible presence of nonlinearity, (often invoked as cause of giant flares in AGN light curves), which is crucial for breaking the degeneracy of models able to reproduce time-averaged spectra and some general timing properties. The result, obtained by utilizing a generalized surrogate method (well suited for unevenly sampled data) and the nonlinear prediction error, indicates the presence of nonlinearity both in the high and, at a slightly lower significance level, in the low count rate states. As a consequence, intrinsically linear models, where the variability is caused by many independent active regions, as magnetic flares or the traditional ``shot noise\" should be ruled out for \\object{Ark~564}. Generalizations of these models, where the spatial and temporal distributions of flares are not random (e.g. Merloni \\& Fabian 2001) might still be a viable solution. The presence of nonlinearity favors nonlinear models as the self-organized criticality model or the emission from the base of a putative X-ray jet, as recently found in several radio-loud AGN. This last hypothesis seems to be supported also by two independent arguments: 1) the need of relativistic beaming effects to explain the extreme values for the efficiency in the conversion of gravitational potential energy into X-ray emission in several NLS1 galaxies (e.g. \\object{PKS~0558-504}, Remillard et al.(1991); \\object{PHL~1092}, Brandt et al. (1999); \\object{RX J1702.5+3247}, Gliozzi et al.(2001)). 2) The recent discovery that, in a sample of 62 NLS1 detected in the FIRST VLA radio survey (Becker \\cite{beck}), $\\sim 40\\%$ of the objects are radio--loud ($R>10$) and the remainder fall in the radio-intermediate range ($1<R<10$) (Whalen et al. 2002).  We then tried to characterize the timing behavior during the low and high count rate states of \\object{Ark~564}. First, we used standard techniques (like the auto--correlation function and excess variance analysis) and linear techniques that are not frequently applied to AGN light curves (like the probability density function), without finding any significant difference between the two states. After showing that, for a limited time interval, the complex light curve of \\object{Ark~564} might be viewed as fractal objects with a non--trivial fractal dimension and statistical self--similarity, we tried to discriminate the two states with the help of the Hurst exponent. Also with this analysis no intrinsic differences were found, leading to the conclusion that the physical processes during the high and low state are of the same nature.  Finally, we have introduced the scaling index spectrum, based on phase space reconstruction, as a tool to discriminate time series. Based on this nonlinear statistic, we showed that the low and high states of \\object{Ark~564}, apparently quite similar with respect to their linear properties, are intrinsically different, with the high state characterized by a higher degree of complexity, meaning that the number of degrees of freedom during the high state is larger that during the low state. However, in order to obtain firmer and more general conclusions we have to await further observations, in particular from the new X--ray missions, with sensitive, high time resolution instruments and the possibility of long continuous observations."
    },
    "0206/astro-ph0206159_arXiv.txt": {
        "abstract": "In this paper, we calculate new line lists and opacities for the 12 bands of the A$^6\\Sigma^{+}$ -- X$^6\\Sigma^{+}$ transitions of the CrH molecule.  Identified in objects of the new L dwarf spectroscopic class (many of which are brown dwarfs), as well as in sunspots, the CrH molecule plays an important role in the diagnosis of low-temperature atmospheres. As a tentative first application of these opacities, we employ our new theoretical CrH data in an atmospheres code to obtain a CrH/H$_2$ number ratio for the skin of the L5 dwarf 2MASSI J1507038-151648 of $\\sim 2-4\\times 10^{-9}$, in rough agreement with chemical equilibrium expectations.  Since in previous compilations the oscillator strength was off by more than an order of magnitude, this agreement represents a modest advance.  However, in order to fit the CrH abundance in the L dwarf spectral class, silicate clouds need to be incorporated into the model.  Given that this subject is still in a primitive stage of development, one should view any spectral model in the L dwarf range as merely tentative.  Nevertheless, a necessary first step in L dwarf modeling is a reliable CrH opacity algorithm and this is what we have here attempted to provide. ",
        "introduction": "\\label{intro}  The CrH molecule had previously been identified in the spectra of sunspots (Engvold, W\\\"ohl, and Brault 1980) and S-type stars (Lindgren and Olofsson 1980), but recently CrH bands have been identified in the spectra of very-low-mass stars and brown dwarfs.  In fact, the 1-0 and 0-0 bands of the A$^6\\Sigma^{+}$ -- X$^6\\Sigma^{+}$ transition of CrH are now used as primary markers for the new L dwarf spectral class (Kirkpatrick et al. 1999a,b). Hence, accurate CrH line lists and oscillator strengths are needed to calculate the CrH opacities now used to model the spectral energy distributions of these transitional and substellar objects. The ground X$^6\\Sigma^{+}$ state of CrH is well characterized from far-infrared (Corkery et al. 1991) and mid-infrared (Lipus, Bachem, and Urban 1991) measurements.  However, until recently the same could not be said for the A$^6\\Sigma^{+}$ state.  In this paper, we describe the recent spectroscopy of CrH and constraining laboratory measurements that now allow us to calculate new opacities due to the A$^6\\Sigma^{+}$ -- X$^6\\Sigma^{+}$ transitions of the CrH molecule. In \\S\\ref{spectrocrh}, we discuss the spectroscopy of the CrH molecule. In \\S\\ref{generation}, we present the procedures and methods by which the new line lists were generated. In \\S\\ref{opac}, we show how we calculate CrH opacities at a given temperature/pressure ($T/P$) point and in \\S\\ref{opacplots} we present representative opacity plots for CrH at temperatures and pressures characteristic of L and brown dwarf atmospheres. In \\S\\ref{bdspec}, we provide and describe a few theoretical spectra generated with the new CrH opacity data and derive an approximate CrH abundance in a representative L5 dwarf atmosphere. Though the focus of this paper is on the derivation of the new CrH line lists and opacities, we use these exploratory synthetic spectra to speculate on various features of observed L dwarf spectra.  We plan to develop our findings in a subsequent paper. ",
        "conclusions": "\\label{conclusions}  Using the vibrational and rotational constants for the A$^6\\Sigma^{+}$ state of the CrH molecule recently derived by Bauschlicher et al. (2001), we have calculated new line lists and opacities for the 12 bands of its A$^6\\Sigma^{+}$ -- X$^6\\Sigma^{+}$ transitions.  The resulting theoretical data are available from the authors. CrH is a defining molecule of the L dwarf spectroscopic class and accurate opacities as a function of temperature and pressure are necessary for spectral syntheses and to extract CrH abundances for L dwarf atmospheres. In a tentative first step, we use the new theoretical data to obtain such an abundance for the L5 dwarf 2MASSI J1507038-151648. The CrH/H$_2$ number ratio we find is $\\sim 2-4\\times 10^{-9}$, in reasonable agreement with expectations."
    },
    "0206/astro-ph0206229_arXiv.txt": {
        "abstract": "{ We report the B band  optical observations of an old (${\\rm\\tau \\sim 17.5}$ Myr)  radiopulsar \\psr \\ obtained with  the Suprime-Cam at the Subaru telescope.  We detected a faint object, B=27$\\fm07\\pm$0.16. Within our  astrometrical accuracy it coincides with the radio position of  the pulsar and with the object detected earlier by Pavlov et al. (\\cite{Pavlov}) in UV with the   \\hst/FOC/F130LP.   The   positional  coincidence   and  spectral properties of the  object suggest that it is the optical counterpart of PSR B0950+08. Its flux in  the  B band is two times higher than one  would expect from  the suggested earlier Rayleigh-Jeans interpretation  of the only available \\hst\\ observations in the adjacent F130LP band.  Based on the B and F130LP photometry of the suggested counterpart and  on the available X-ray data  we argue in favour of nonthermal origin  of the broad-band optical  spectrum of  \\psr, as it is observed for the optical emission of the younger, middle-aged pulsars PSR  B0656+14 and  Geminga.  At the  same time, the optical efficiency of  \\psr, estimated from its spin-down power and the detected optical flux, is  by several orders  of magnitude higher  than for these pulsars, and  comparable with that for the  much younger and more energetic Crab  pulsar. We cannot  exclude the presence of  a compact, $\\sim$  1\\arcsec,  faint pulsar  nebula  around  \\psr,\\  elongated perpendicular to the  vector of its proper motion, unless  it is not a projection  of  a  faint  extended  object on the  pulsar  position. ",
        "introduction": "PSR  B0950+08  is the fourth in  the  list of the first radio  pulsars discovered in 1968  (e.g., \\cite{Bell}).  Since that time  it has been extensively studied  in a  wide radio frequency  range, from  0.102 to 10.55 GHz.  The pulsar parameters (period $P$ and its derivative $\\dot P$,  age  $\\tau$,  magnetic  field  $B$,  dispersion  measure  $D\\!M$, spin-down  luminosity  $\\dot  E$,  distance  $d$,  position  $\\alpha$, $\\delta$,  proper  motion  $\\mu_\\alpha$,  $\\mu_\\delta$,  and  parallax $\\pi$) are shown  in Table \\ref{t:par}. So far this is the oldest pulsar  among ordinary  pulsars detected outside  the radio range.   \\begin{table*}[t] \\caption{ Parameters of   \\psr\\ (from \\cite{Taylor}, unless specified otherwise). } \\begin{tabular}{ccccccccccc} \\hline\\hline \\multicolumn{5}{c}{Observed}&&\\multicolumn{5}{c}{Derived} \\widerul\\\\ \\cline{1-6}\\cline{8-11} $P$\t& $\\dot P$   & $D\\!M$       & $\\alpha_{2000}$, $\\delta_{2000}^{a,b}$                             & $\\mu_\\alpha$, $\\mu_\\delta^b$     & $\\pi^b$  && $\\tau$ & $B$          & $\\dot E$ & $d^{b,c}$     \\widerul \\\\ ms      & $10^{-15}$ & cm$^{-3}$ pc &                                                                    & mas yr$^{-1}$                    & mas    && Myr    & G  & erg $s^{-1}$ & pc       \\\\ \\hline 253     & 0.229      & 2.97         & \\twolines{$9^h53^m09\\fs313(3)^d$}{7$^\\circ$55\\arcmin36\\farcs08(4)} & \\twolines{$-$2.09(8)}{29.46(7)}  & 3.82(7) && 17.5   & $0.24\\times10^{12}$         & $5.6 \\times 10^{32}$   & $262(5)$ \\widerul\\\\ \\hline \\end{tabular} \\begin{tabular}{ll} $^a$\\ The  position at the epoch of the Subaru observations, MJD 51930. % & $^d$\\ The numbers in parentheses are uncertainties \\widerul\\\\ $^b$\\ See Brisken et al.  (\\cite{Brisken2}) &\\phantom{$^d$}\\ referring to the last significant digit quoted, \\\\ $^c$\\ Parallax based distance &\\phantom{$^d$}\\ e.g., $1.234(56)=1.234\\pm 0.056$ \\\\ \\end{tabular} \\label{t:par} \\end{table*}     The X-ray counterpart of PSR B0950+08 was first found by \\cite{Seward} and  \\cite{Cordova}  with  the  Einstein observatory.  Later,  it  was observed with  the \\rosat\\ by Manning and Willmore  (\\cite{Manning}), who analyzed  the soft X-ray spectrum  in  (0.08--2.4)  keV  range. The  blackbody  (hereafter  BB) spectral fit  gave a  temperature $T_{\\rm BB}  = (2.1 \\pm  0.6) \\times 10^6  $~K (90\\%  confidence limits)  and a  very small  radius  of the emitting  region  $R_{\\rm BB}  \\le  20$~m.   The  pulsar distance  was assumed  to  be  of  about  130  pc  based  on  early  radio  parallax measurements by \\cite{Gwinn} with  the VLBI.  The power law (hereafter PL, $F_\\nu\\propto\\nu^{-\\alpha_\\nu}$) spectral fit was also acceptable, but, owing  to the poor count statistics  for this faint  X-ray object, it resulted  in   a  very  uncertain   spectral  index  ${\\rm   -0.2  \\le \\alpha_{\\nu} \\le 3.1}$.  Some evidence, at 99\\% significance level, of the pulsed X-ray  emission with pulsed fraction ${\\rm  0.34 \\pm 0.18}$ was  found  by Wang and Halpern (\\cite{Wang})  using  archival  \\asca\\  observations  in (0.5--5) keV range.   The \\asca\\ spectral data were  also equally well fitted  either by  the BB  with a temperature  $T_{\\rm  BB} = (5.70\\pm0.63)\\times 10^6 $~K, higher than that in the \\rosat\\ range, and with $R_{\\rm  BB}   \\sim  14$~m,  or   by  the  PL  with   $\\alpha_{\\nu}  = 0.40\\pm0.22$.  In accordance with NS cooling theories, the temperature of the  whole surface of  this relatively old  NS should be  less than $10^5$~K (e.g.,  \\cite{Yakovlev}). This is much less  than observed by both the \\rosat\\ and \\asca.  The  high temperature and small emitting area inferred from the BB fit can be explained by the presence of small hot polar caps  at the NS surface  produced by the  impact of relativistic particles from  the pulsar  magnetosphere, see, e.g.,  Wang and Halpern (\\cite{Wang}) and references therein.    An optical counterpart of  PSR B0950+08 was suggested by Pavlov et al. (\\cite{Pavlov}) based  on observations of the  pulsar field with  the \\hst/FOC with the  long-pass F130LP  filter ($\\lambda\\lambda  = 2310-4530$  \\AA).  A faint ($m_{\\rm F130LP}=27$\\fm$1$) point-like object was found with the 1\\farcs85 offset from the pulsar radio position.  The offset was later revised and decreased to $\\simeq 1$\\arcsec\\   by Pavlov (\\cite{Pavlov1}). If this is a pulsar, it  is the  faintest pulsar ever detected in the optical. For comparison, the visual magnitude  of the young  Crab pulsar,  which is  about ten  times more distant but  $1.75\\times 10^4$ times younger, is  16$\\fm65$ (e.g., see the review  by \\cite{Mignani2}).  Assuming that the  detected optical object is the  pulsar, Pavlov et al. (\\cite{Pavlov}) showed that the extension of the \\rosat\\ and  \\asca\\ BB  fits into the  UV-optical range gives  a flux lower than observed by  3--4 orders of magnitude. This excludes thermal radiation from  the pulsar polar caps as  a possible source of the optical  radiation. The assumption that the  detected optical flux is due to thermal  emission from the entire surface of a  NS with a BB radius $R_{\\rm BB}=13$~km yielded  the surface temperature $T_{\\rm BB} \\sim 7 \\times 10^4$~K at ${\\rm  d=130}$~pc. This is still too high for the $\\sim 1.75  \\times10^7$~yr cooling NS and can  be only explained by some   reheating   mechanism   operating   inside  the   star   (e.g., \\cite{Miralles}).  On  the other  hand, the optical  flux would  be in agreement with  the extension of the  PL X-ray fit if  the index ${\\rm \\alpha_{\\nu}}$ lies  within the 0.26--0.35 range.  It  is in agreement with  the \\rosat\\  and \\asca\\  data but  needs a confirmation  by deeper observations in X-rays and by  the detection of the counterpart in other optical bands.   Possible  detection of  PSR  B0950+08 in  the R  band with  the 6m telescope BTA  has been  reported by \\cite{Sokolov}  and \\cite{Kurt1}. An  object  with  $R=25\\fm4\\pm0.3$  was marginally  ($S/N  \\simeq 3$) detected  with  a  short  exposure  in  poor  seeing  conditions.   If confirmed,  this  detection  suggests  that  the pulsar  may  be  much brighter in  the optical  and may  have a very  steep increase  of the spectrum towards  longer wavelengths  than one would expect  from the detection in the near UV.  In this paper, we report  the observations of the PSR~B0950+08 field in the B  band with the Subaru  telescope.  We analyze  our data together with the available  optical-UV data from the \\hst,  and with the X-ray data  from the \\asca\\  and \\rosat,\\  making use  of the  recent  much more accurate radio measurements of the pulsar proper motion, parallax, and distance  with  the   VLBA  by Brisken et al. (\\cite{Brisken2}). Observations and data reduction are  described in Sec.~2. In Sec.~3 we present the astrometrical referencing and photometry, and in Sec.~4 we discuss the results and their implications.   \\begin{table}[b] \\caption{Log of observations of the pulsar in the B band.} \\begin{tabular}{cccccc} \\hline\\hline Exposure\t&       UT  \t& Duration\t   & Airmass & Seeing    \\\\ number\t\t& 21 Jan 2001\t&   sec            &         &  arcsec   \\\\ \\hline 1\t\t& 12:41     \t&   600            & 1.037   &  0.65     \\\\ 3\t\t& 13:21     \t&   600            & 1.085   &  0.66     \\\\ 4\t\t& 13:54     \t&   600            & 1.129   &  0.69     \\\\ 5\t\t& 14:08     \t&   600            & 1.164   &  0.72     \\\\ 6\t\t& 14:22     \t&   600            & 1.205   &  0.73     \\\\ 8\t\t& 14:50     \t&   600            & 1.313   &  0.75     \\\\ \\hline \\end{tabular} \\label{t:log} \\end{table} ",
        "conclusions": "\\label{s:disc}  In Fig.~\\ref{f:flux} we combine our results with the available data from the optical and X-ray observations of \\psr. In the optical range we neglect the interstellar extinction, which is expected  to be very low for this pulsar, $E_{B-V}\\la 0.02$ (e.g., Pavlov et al. \\cite{Pavlov}). We excluded from our consideration the observations in the R band by \\cite{Kurt1} because of very low signal-to-noise ratio. In the X-ray range we show unabsorbed fluxes resulted  from  the BB and PL spectral fits of the \\rosat\\  (Manning and Willmore \\cite{Manning}) and \\asca\\ (Wang and Halpern \\cite{Wang}) data. The fits are rescaled to the new distance value $d=262\\pm5$ pc  measured by Brisken et al.  (\\cite{Brisken2}).  Within the errors the flux $F_{\\rm B}$ coincides with the value $F_{\\rm F130LP}=0.051\\pm0.003$~$\\mu$Jy measured by  Pavlov et al. (\\cite{Pavlov}) with the \\hst/FOC in the F130LP band. The B band ($\\lambda\\lambda \\sim 3944-4952$~\\AA) considerably overlaps with the F130LP band ($\\lambda\\lambda \\sim 2310-4530$~\\AA). However, their  pivot  wavelengths $\\lambda^0_{\\rm B}=4448$~\\AA\\  and $\\lambda^0_{\\rm F130LP}=3750$~\\AA\\ are different. This difference and close flux values in two bands suggest a flat spectrum of the object in a wide  $\\lambda\\lambda \\sim 2300-5000$~\\AA\\ spectral range. Along with the positional coincidence, we consider this  flat spectrum, which is typical for pulsar optical spectra (e.g., Koptsevich et al. \\cite{Koptsevich}; Mignani \\& Caraveo \\cite{Mignani3} ), as an additional argument in favour of the detection of the \\psr\\ optical counterpart.  If the detected object is the pulsar,  its flux in the B band does not follow the  Rayleigh-Jeans law  suggested by  Pavlov et al. (\\cite{Pavlov})  to explain the optical-UV radiation in the  only F130LP band as a low temperature thermal  emission from the  entire surface  of a  reheating/cooling NS with $R_{\\infty}=13$ km.   In that case in the B  band we would  detect twice smaller flux  than  measured.  This is well outside the  uncertainties and  suggests  a nonthermal  origin  of the  pulsar emission at least in the B band.   \\begin{table*} \\caption{ Parameters of the radio pulsars and the efficiencies of their optical emission $\\eta_B$ in the B band. } \\begin{tabular}{lccccccc} \\hline\\hline Source  \t& $\\log\\tau$\t& $d$    \t\t\t\t\t& $\\log\\dot{E}$\t& $B$ \t\t& $M_{\\rm B}$\t\t& $\\log L_{\\rm B}$\t\t& $\\log\\eta_{\\rm B}$ \\\\ &  yr\t\t& kpc      \t\t\t\t\t& erg s$^{-1}$\t& mag \t\t& mag \t\t\t& erg s$^{-1}$\t\t\t&           \\\\ \\hline Crab\t\t&  3.1       \t& 2.0(1)           \t\t\t\t& 38.65\t\t& $15.25(7)^1$\t& 3.74(13)   \t\t& 33.23(5)\t\t\t& $-$5.42(5)   \\\\ B0540$-$69\t&  3.2       \t& $50^{+5}_{-0.6}$          \t\t\t& 38.17\t\t& $22.0(3)^2$\t& 3.48(37)\t\t& 33.47(15)\t\t\t& $-$4.7(2) \\\\ Vela\t\t&  4.1       \t& ${0.294^{+0.076}_{-0.050}}^a$\t\t\t& 36.84\t\t& $23.7^(3)^3$\t& $16.4^{+0.5}_{-0.8}$\t& 28.3(3)\t\t\t& $-$8.5(3)   \\\\ B0656$+$14\t&  5.0       \t& ${0.5^{+0.26}_{-0.3}}^b$           \t& 34.58\t\t& $25.15(13)^4$\t& $16.8^{+2.1}_{-1.0}$\t& $28.2^{+0.4}_{-0.9}$\t\t& $-6.4^{+0.4}_{-0.9}$\\\\ Geminga\t\t&  5.5       \t& ${0.153^{+0.059}_{-0.034}}^{c}$ \t\t& 34.51\t\t& $25.7(3)^5$\t& $19.8^{+0.8}_{-1.0}$\t& $26.95^{+0.16}_{-0.10}$\t& $-7.56^{+0.16}_{-0.10}$  \\\\ B1929$+$10\t&  6.5       \t& $0.331(10)^{d}$ \t& 33.59\t\t& $ \\geq26.2^6$\t& ${20.0^{+0.2}_{-0.2}}^{**}$\t\t& ${27.26^{+0.2}_{-0.3}}^{**}$\t\t& ${-6.3^{+0.2}_{-0.3}}^{**}$ \\\\ B0950$+$08\t&  7.2      \t& 0.262(5)$^{d}$\t\t\t\t& 32.75\t\t& $27.07(16)^7$\t& 19.98(19)\t\t& 26.88(8)\t\t\t& $-$5.87(8)  \\\\ \\hline \\hline \\\\ \\end{tabular} \\label{t:eff} \\begin{tabular}{l} The B band magnitudes are from: \\\\ $^1$ \\cite{Percival}; $^2$ \\cite{Middleditch} (spectroscopic data by \\cite{Hill} give $\\sim 1^{m}.5$ smaller magnitude); \\\\ $^3$ \\cite{Nasuti}; $^4$ Koptsevich et al. \\cite{Koptsevich}; $^5$ \\cite{Bignami}; $^6$  Pavlov et al. (\\cite{Pavlov}) (3$\\sigma$ upper limit); \\\\ $^7$ This work. The distances $d$ are from : \\\\ $^a$ \\cite{Caraveo}; $^b$ \\cite{Taylor}, \\cite{Finley}, \\cite{Anderson}; $^c$ \\cite{Caraveo2} \\\\ $^d$ Brisken et al.  (\\cite{Brisken2}); \\\\ $^{**}$ The estimates are based on the counterpart detection in the adjacent F130LP band by Pavlov et al.  (\\cite{Pavlov}), see Sec.~\\ref{s:disc}.\\\\ Uncertainties of $M_B$, $L_B $ and $\\eta_B$  include uncertainties of the optical flux and distance measurements. \\end{tabular} \\end{table*}    If the flux in the F130LP band is still dominated by  the Rayleigh-Jeans tail, we would obtain $T_{\\rm BB}=(3.4\\pm0.3)\\times 10^5$  K, which is a factor of four-five higher than found by  Pavlov et al. (\\cite{Pavlov}) mainly owing to the change of the  distance to  \\psr \\ from 130~pc to 262~pc (see the stripe-filled BB belt crossing the F130LP band and extended to X-rays in Fig.~\\ref{f:flux}). A NS with such a hot surface would  produce a bolometric flux $ \\sim (0.8-1.9)\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ in the ($0.08-2.4$) keV  band (see a big dashed cross in  Fig.~\\ref{f:flux}). It is well above the value $(2.4 - 7.3)\\times 10^{-14}$  erg cm$^{-2}$ s$^{-1}$, measured by Manning and Willmore (\\cite{Manning}) with the \\rosat\\ under the assumption of the BB spectrum of the detected X-ray radiation (marked by a big cross below the dashed one in Fig.~\\ref{f:flux}). This means that the whole surface of \\psr\\ is actually much cooler and its emission cannot dominate  in the F130LP band. Thermal emission from hotter, $T_{\\rm BB}=(5.7 \\pm0.6) \\times 10^6$~K, but much smaller polar caps of the pulsar with $R_{\\rm pc} \\sim 28.6$~m, which may explain the detected X-ray radiation (Wang and Halpern \\cite{Wang}), can  hardly be  visible in the optical range also because of very small areas of the caps  inferred from the BB fit of the X-ray data (see square-filled BB belt in Fig.~\\ref{f:flux}).   For the above reasons it is most likely that the optical radiation of  \\psr \\ is completely dominated by nonthermal emission produced in the magnetosphere of the rapidly rotating NS, as it is believed to be for young and well studied pulsars like the Crab and PSR  B0540$-$69. Within large uncertainties of the available X-ray data for \\psr \\ the PL with the spectral index $\\alpha_{\\nu} \\simeq 0.18$ matches both the X-ray and optical fluxes  including the B and F130LP bands.  This value is consistent with what was obtained by Wang and Halpern (\\cite{Wang}) from the  analysis of the  \\asca\\ X-ray data, $F_{\\rm ASCA}^{\\rm PL} = 2.77~(+1.33,-1.05)\\times 10^{-13}$ erg cm$^{-2}$ s$^{-1}$, ~$\\alpha_{\\nu}=0.40 \\pm 0.22$ in (0.5 -- 5) keV energy range (see the ASCA cross and the brick-filled region in Fig.~\\ref{f:flux}). The PL fit of the X-ray data is more preferable since the BB fit, which is statistically also acceptable,  implies by an order of magnitude smaller emitting area than it is predicted by standard models of hot polar caps at the pulsar surface (e.g., \\cite{Arons} ). The inferred spectral index differs from $\\alpha_{\\nu} \\sim 1$ which is typical for nonthermal soft X-ray radiation of  most rotation-powered NSs (e.g., \\cite{Becker}). However, it can be as low as 0.4 for middle-aged pulsars (Koptsevich et al. \\cite{Koptsevich}) and we cannot exclude  a decrease of the slope of the \\psr \\  spectrum  towards the optical range as it is seen in the case  of the Crab pulsar (e.g., Crusius-W\\\"{a}tzel et al. \\cite{Crusius}). More X-ray and optical data are needed to check the spectral shape  for \\psr. With  new data it would be also useful to perform  the BB $+$ PL, and/or NS atmosphere $+$ PL  fits (e.g., Zavlin et al. \\cite{Zavlin}) to better  constrain the parameters of the nonthermal and thermal spectral components from the pulsar magnetosphere and polar caps and to estimate their contribution to the pulsar emission in different spectral bands.   \\begin{figure}[t] \\begin{center} \\includegraphics[width=88mm,clip]{figure5.eps} \\end{center} \\caption{The  optical luminosity  $L_{\\rm B}$  ({\\it  top}), spin-down power $\\dot{E}$  ({\\it middle}) and  optical efficiency in the  B band $\\eta_B$ ({\\it bottom}) for all  radio pulsars detected in the optical range as functions of pulsar age $\\tau$ (see also Table~\\ref{t:eff}).} \\label{f:eff} \\end{figure}  If the detected object is \\psr, the pulsar optical luminosity in the B band, assuming  isotropy   of  the   emission,  is  $L_{\\rm   B}  =   4  \\pi d^2F_\\nu\\Delta\\nu = (7.6\\pm1.2)\\times10^{26} d^2_{262}$~ erg s$^{-1}$, where $  d_{262}=d/(262$~pc) is the normalized distance.  As seen from Table~\\ref{t:eff} and Fig.~\\ref{f:eff} (top panel), where we collected the data in the B band  for all pulsars detected in the optical range, $L_{\\rm  B}$ of  \\psr\\ is  comparable with  the luminosity  of the middle-aged  Geminga pulsar  but  by almost  six  orders of  magnitude smaller than $L_{\\rm  B}$ of much younger and energetic Crab  pulsar.  These data indicate  that, starting approximately  from the  age of  Geminga, the decrease of the optical luminosity  with the pulsar age becomes slower than  at early  stages  of  the pulsar  evolution  or even  disappears completely, while the spin-down  luminosity continues to decrease with approximately the same rate (Fig.~\\ref{f:eff}, middle panel). This means that for unknown reasons mechanisms  of the optical emission may start to be more efficient again after $\\sim 10^6$~yr age.  The ratio of the optical  luminosity  to the  pulsar  spin-down luminosity,  $\\eta_{\\rm B}=L_{\\rm   B}/\\dot{E}$,  which   can   be  considered   as  a   rough characteristic  of  this efficiency,  is  plotted in  Fig.~\\ref{f:eff} (bottom panel).  It  is seen that optical photons  are produced by the old  PSR B0950+08 with  almost the  same efficiency  as by the young Crab pulsar.   In  this respect  it  is not  surprising  that  we see  some indications  of the  presence of  a compact  pulsar nebula  around \\psr. Manning and Willmore \\cite{Manning} discuss a compact synchrotron nebula,  which is unresolved with the \\rosat, as the most likely explanation of the observed X-ray spectrum of the pulsar. The estimated sizes of the possible X-ray nebula $\\sim (6\\arcsec - 1\\arcmin)$ are much larger than $\\sim 1\\arcsec$  we  see in our optical images. But the difference in the X-ray and optical sizes is typical for known compact pulsar  nebulae where only internal brightest structures of a nebula are visible in the optical range (e.g., Weisskopf et al.~\\cite{We}).  The efficiency derived by  Pavlov et al. (\\cite{Pavlov}) in the F130LP band for another old, $\\sim 3.1\\times 10^{6}$ yr, PSR B1929+10 is about $ (3-8) \\times 10^{-7}d^2_{170}$. This pulsar has not yet been detected in the adjacent B band, but we can assume that its flux in B is close to that in F130LP,  as it is  for   PSR  B0950+08.  Scaling by the  factor $ \\Delta  \\nu_{\\rm B}/\\Delta  \\nu_{\\rm F130LP}  = 0.24$  and by  the new radio parallax based distance  $d=331$~pc (Brisken et al. \\cite{Brisken2}) yields $\\eta_{\\rm  B} \\sim  (2.7-7.3) \\times  10^{-7}d^2_{331}$. With this value PSR B1929+10 occupies an intermediate  position at the rising part of the dependence of $\\eta_{\\rm B} (\\tau)$, between Geminga and \\psr,\\ as it is expected from its age. Although our consideration of the efficiency evolution is based  on the data obtained  only in the B band, it appears to be qualitatively valid for the whole optical range since the  broad-band spectra of  all pulsars are  almost flat. Hence, the differences  of the bolometric  optical fluxes of  various pulsars should not be affected  strongly by insignificant differences of their spectral slopes in this range.   Physical reasons for such a high increase of the optical efficiency at late stages of the pulsar evolution  are not quite clear.  We can only note   that   the   efficiency   for gamma-ray radiation, $\\eta_{\\gamma}=L_{\\gamma}/\\dot{E}$,  also   appears  to  be  generally higher for older pulsars  (e.g., \\cite{Thompson}). It is difficult to estimate statistical significances of these facts and their possible correlation  since the numbers of pulsars currently detected in  the optical and gamma-rays are too small (about ten only). Morever, not all known gamma-ray pulsars are detected in the optical range (and {\\it vice versa}). However, the observed tendency in both  ranges seems to be interesting and can  hardly be ignored. It is obvious that in the both wavelength  ranges  the radiation is  nonthermal  and   originates  in magnetospheres  of rapidly rotating  NSs. Thus,  the increase  of the efficiencies  in two  very  different ranges  may  reflect an  overall increase  of the  magnetospheric  activity with  the  NS spin-down.  A global  electrodynamic  model of  the  pulsar  magnetosphere with  the activity caused by the magnetic field-aligned potential drop producing electron-positron  pairs   in  the  magnetic  polar   regions  of  the magnetosphere predicts  that the  efficiency should increase  for high energy photons $ \\propto  P^2$ (\\cite{Shibata}).  However, the optical data, particularly  for young pulsars,  do not follow  this dependence and detailed  studies of  the electrodynamics and  radiation processes still have to be done to explain the efficiency evolution in different spectral  ranges.  Further  observations of  the candidates to the optical counterparts of the old pulsars PSR B0950+08 and  PSR B1929+10  in different  spectral bands  would be  very  useful to resolve the efficiency problem and to better understanding the nature of the optical emission of  pulsars and  the evolution of  this emission  with pulsar age.  Measurements  of  their  proper  motion  and  the  detection  of pulsations with the pulsar periods in the optical range would be most important to provide firm  evidence   of  the  pulsar   nature  of  the   detected  optical objects. Simultaneous  studies of the optical and  X-ray pulse profile would  provide  stronger indications  whether  the  optical and  X-ray emissions are generated by the same physical process.  Our observations of a very faint  \\psr \\ with the  Subaru show that new  generation of large ground-based telescopes is very  effective for these studies and could lead to a  considerable increase of the number  of pulsars detected in the optical range in the near future."
    },
    "0206/astro-ph0206184_arXiv.txt": {
        "abstract": "We report the first time-series observations of the short outburst of the proposed intermediate polar HT Cam (=RX J0757.0+6306). On 2001 December 29, we detected the object was undergoing a bright outburst at the magnitude of $m_{vis}=12.2$. Following this detection, we started international joint observations through VSNET. The light curve showed a gradual decline for the first 0.5~d. Following this short plateau phase, the rate of decline dramatically increased to more than 4~mag d$^{-1}$. Within 1.5~d from the outburst detection, the object almost declined to the quiescent level. During the rapidly declining phase, long-term modulations with a period of 86~min and strong pulses with a period of 8.6~min were observed. We concluded that 86~min and 8.6~min are the orbital period and the spin period of HT Cam, respectively. By the detection of the spin period, we confirmed the IP classification of HT Cam. However, its outburst behavior rather resembles that of dwarf novae. The discrepancy between the declining rates of the total flux and the pulse flux strongly suggests that the disk instabilities were taking place during the outburst. ",
        "introduction": "Cataclysmic variables consist of a white dwarf primary and a late-type main-sequence companion star which is mass-losing by the Roche-lobe overflow. The white dwarf accretes the mass from the secondary through an accretion disk in non-magnetic systems. In magnetic systems, however, either the formation of the disk is fully prevented, or the inner disk is truncated, depending on the magnetic fields of the white dwarf and the binary parameters. Within the magnetosphere radius the mass falls along the magnetic field toward poles of the white dwarf, where radiations from X-ray to optical emission is originated. Magnetic systems are divided into two classes of polars and intermediate polars (IPs). The polars are the systems with strong fields and a synchronously rotating white dwarf, and are soft X-ray sources. IPs are ones with weaker fields and an asynchronously rotating white dwarf, and are hard X-ray sources. Reviews for the magnetic systems are given by \\citet{cro90polar}, \\citet{pat94ipreview}, \\citet{war95book}, and \\citet{hel96IPreview}.  There are six IPs in which outbursts are observed: GK Per, EX Hya, XY Ari, DO Dra\\footnote{There has been some confusion regarding the nomenclature of this object. Some authors call the same variable as YY Dra. We consistently use the name DO Dra, following the original nomenclature by the General Catalogue of Variable Stars (GCVS) \\citep{NameList67,kho88dodra}.}, TV Col, and V1223 Sgr. GK Per shows long outbursts that last for 2~months \\citep{sab83gkper,kim92gkper}. Remaining 5 IPs show rather short outbursts. Outbursts of XY Ari \\citep{hel97xyarioutburst} and DO Dra \\citep{wen83dodra,szk02dodra} are with amplitudes of $3-5$~mag and durations of $\\sim 5$~d. TV Col \\citep{szk84tvcolflare,hel93tvcolperiods} and V1223 Sgr \\citep{vaname89v1223sgr} show very short low-amplitude outbursts that last only $\\sim 0.5$~d. Outbursts of EX Hya have intermediate  characteristics between the previous two types, that is, with amplitudes of $\\sim 3.5$~mag and durations of $2-3$~d \\citep{hel89exhya}.  Outbursts of GK Per are explained as inside-out outbursts by thermal instability in the disk, though the atypically large inner radius and mass transfer rate are required to fit the observed recurrence time and duration of outbursts \\citep{kim92gkper}. Outbursts of XY Ari and DO Dra rather resemble outbursts of dwarf novae and also explained as inside-out outbursts by disk instability \\citep{hel97xyarioutburst}. The recurrence time one order longer than the typical value is thought to be due to inner truncation of the disk \\citep{ang89DNoutburstmagnetic}. On the other hand, the short-lived and low-amplitude outbursts of TV Col and V1223 Sgr are thought to be mass transfer outbursts \\citep{hel93tvcol}. \\citet{hel00exhyaoutburst} suggested that outbursts of EX Hya are also mass transfer events.  HT Cam is a cataclysmic variable identified as the optical counterpart of the hard X-ray source RX~J0757.0$+$6306, discovered during the ROSAT All-Sky Survey. Tovmassian et al. (\\citeyear{tov97htcam},\\citeyear{tov98htcam}) suggested that this object is an intermediate polar with a shortest orbital period of 80.92 min and a spin period of 8.52 min. However, the existence of dwarf nova-like outbursts and the short orbital period allowed an alternative interpretation that it may be an SU UMa-type dwarf nova or WZ Sge-type stars \\citep{tov98htcam}. After the discovery, five outbursts were recorded (vsnet-alert 1379, 3025, 4845, 5255, 6066, and \\cite{wat98htcam}). All of them were very short and only one or a few point observations were obtained for each one.  Here we report the first time-series observations during its short outburst. ",
        "conclusions": "\\subsection{8.6 min oscillation}  Short-term coherent pluses with a period of 8.6~min were observed during the outburst. This period is thought to be the spin period of white dwarf. \\citet{tov98htcam} found a period of 8.5~min from their quiescence observations. We could not detect the 8.6~min period after the object declined except on Jan. 3. The pulse amplitude should have been below the limit of detection, $\\sim 0.2$~mag.  The pulse amplitude in flux scale continuously reduced as the system declined, but in magnitude scale it decreased once and increased again during the rapidly declining phase (figure \\ref{fig:spin_profile}), that is, the declining rate of the pulse flux was slower than that of the total flux of the system. The pulse flux and the total flux roughly correspond to the accretion rate onto the white dwarf and the accretion rate through the outer region. The discrepancy of declining speeds of two accretion rates suggests the existence of a disk, rather than via ballistic accretion.  \\subsection{86 min oscillation}  \\citet{tov98htcam} suggested 81~min as the orbital period of HT Cam from the radial velocity curves, but they could not reject 76~min and 86~min. Their quiescent photometry showed no modulations with the period of $\\sim$81~min. This indicates the low inclination of HT Cam. In our observations, after the declining to quiescence, no modulations with the orbital period were observed, however, large-amplitude quasi-periodic modulations with the period of 86~min were observed during the declining phase of the outburst. These modulations are likely related to the orbital period.  If the orbital period of HT Cam is 81~min, our period is longer than it. During superoutbursts in SU UMa-type dwarf novae, modulations with a period a few percent longer than its orbital period, called superhumps are observed \\citep{war85suuma}. It is known that superhumps are detected regardless of the inclination of the system. However, in conjunction with the short duration of the outburst, we consider the 86-min unlikely represents the superhump period.  This period more likely reflects the true orbital period among the aliases listed in \\citet{tov98htcam}.  In WZ Sge stars, a small sub-group of SU UMa-type dwarf novae with very short orbital periods of around 80--90~min, modulations with orbital periods, called early superhumps, appear in the early stage of superoutbursts \\citep{mat98egcnc}. Early superhumps have double-humped plofile regardless of the inclination of the system. However, the averaged profile shown in figure \\ref{fig:1st_night} is not double-humped. Thus the possibility that the modulations were early superhumps is excluded.  The existence of the orbital modulations only during the outburst may be due to the changing aspect of a irradiated surface of the secondary star, an enhanced hot spot moving into and out of our line of sight, partial eclipse of the thickend and brightened disk, or a combination of them.  \\subsection{Long-term variability}  The visual observations covering four years with fairly good density have detected 6 rather regular outbursts with an amplitude of $\\sim 4$~mag and a duration of $< 3$~d. Our observations on the 6th outburst of HT Cam confirmed the short duration and rapid fading, perhaps slightly slower than rising. The outburst amplitude was $\\sim 4.5$~mag, though its duration was very short.  EX Hya shows outbursts with a similar short duration of $\\sim 3$~d and the amplitude of $\\sim 3$~mag. \\citet{hel00exhyaoutburst} mentioned that the outbursts of EX Hya are mass transfer events rather than disk instabilities, due to the rarity and irregularity of the outbursts (e.g. no outburst was observed in a period of more than 10 yr, however, in another period, two outbursts occurred with an interval of only 8~d), the declining rate, the increased equivalent widths of emission lines and the enhanced mass-transfer stream during the outbursts.  However, the outburst profile, which shows rather slower declining at first and then very fast declining later, and the regularity of the outburst recurrence time of HT Cam are similar to those of DO Dra. DO Dra shows outbursts with a duration of $\\sim 5$~days and a recurrence time of 868~d \\citep{sim00dodra}. They are explained by disk instability as dwarf novae \\citep{szk02dodra}.  The differences of HT Cam from DO Dra are the shorter outburst duration within $3$~d, and the shorter cycle length of $\\sim 150$~d. The recurrence time of HT Cam is quite typical value among the dwarf novae, compared with those of two intermediate polars whose outbursts are likely explained by disk instabilities ( $\\sim 870$~d of DO Dra as written above, and $\\sim 900$~d or more of GK Per; \\cite{sim02gkper}). There are several dwarf novae that have recurrence times about or longer than $\\sim 150$~d and show short outbursts about 3--5~d (e.g. HT Cas, EK TrA, IR Com).  \\citet{ang89DNoutburstmagnetic} calculated the effects of the inner disk truncation on the outbursts in the intermediate polars: the duration becomes shorter in both cases of mass transfer instability and disk instabilities, and for the outbursts by disk instabilities, the interval becomes longer if the outburst starts at the inside, however does not change if the outburst starts at the outer disk edge.  Using this result, we can explain the outburst behavior of HT Cam as follows. The outbursts are caused by disk instabilities starting at the outside, thus the recurrence time is comparable to those of dwarf novae, $\\sim 150$~d. The short duration and the rapid decline of the outburst are due to the inner truncation of the disk. This is consistent with the rather short spin period indicating intermediate magnetic fields \\citep{kin91IPspinevolution}, and the discrepancy between the declining rates of pulse flux and total flux strongly suggesting the existence of the disk."
    },
    "0206/astro-ph0206467_arXiv.txt": {
        "abstract": "Broad Fe~K$\\alpha$ emission lines have recently been reported in a number of Galactic black holes.  Such lines are useful accretion flow diagnostics because they may be produced at the inner accretion disk and shaped by relativistic effects, but in general they have only been observed at luminosities of $L_{X} \\sim 10^{37-38}~{\\rm erg}~{\\rm s}^{-1}$ in soft X-rays.  The Galactic microquasar V4641~Sgr --- widely known for its 12.2~Crab (1.5--12 keV) outburst in 1999 September --- displayed low-level activity in 1999 March. \\textit{BeppoSAX} observed the source in this state and Fe~K$\\alpha$ line emission was found (in~'t Zand et al. 2000).  In re-analyzing these data, we find strong evidence that the Fe~K$\\alpha$ line profile is broadened.  For the most likely values of the source distance and black hole mass measured by Orosz et al. (2001), our fits to the total spectrum indicate that the source was observed at a luminosity of $L_{X} = 1.9^{+1.0}_{-0.8} \\times 10^{36}~{\\rm erg}~{\\rm s}^{-1}$ (2--10 keV), or $L_{X}/L_{Eddington} = 1.8^{+0.9}_{-0.8} \\times 10^{-3}$.  Advection-dominated accretion flow (ADAF) models predict a radially-recessed disk in this regime.  In contrast, fits to the observed Fe~K$\\alpha$ emission line profile with a relativistic line model constrain the inner disk to be consistent with the marginally stable circular orbit of a Schwarzschild black hole. ",
        "introduction": "Recently, studies of the Fe~K$\\alpha$ line profiles in Galactic black holes and black hole candidates have revealed strong, skewed lines, similar to those that are observed in some Seyfert galaxies (see, e.g., Fabian et al. 2000; Weaver, Gelbord, \\& Yaqoob 2001).  Such lines are likely produced by the irradiation of an accretion disk which extends close to the marginally stable orbit around the black hole.  The exact shape of such lines can indicate the inner extent of the disk (and thereby constrain black hole spin), its ionization state, and the nature of the interaction between the disk and the corona that irradiates its surface.  The broad Fe~K$\\alpha$ lines recently reported have all been observed in bright persistent systems (Cygnus~X-1: Miller et al. 2002a, GRS~1915$+$105: Martocchia et al. 2002) or in bright states of transient systems (XTE~J1650$-$500: Miller et al. 2002b, SAX~J1711.6$-$3808: in~'t~Zand et al. 2002).  These lines provide important new constraints on the black hole spin and inner accretion flows of these systems.  While most models for the accretion flow geometry in Galactic black holes predict that the disk extends near to the marginally stable orbit in bright states ($L_{X}\\simeq 10^{37-38}$~erg/s, or higher), the nature of lower-luminosity states may be different.  In particular, ADAF models for transient systems (e.g., Esin, McClintock, \\& Narayan 1997) predict that the disk may be recessed to several tens or hundreds of gravitational radii ($R_{g} = G M_{BH} / c^{2}$).  V4641~Sgr is usually identified with its remarkable 12.2~Crab outburst lasting only a few hours on 1999 September 15 (Smith et al. 1999). Dynamical constraints strongly indicate a black hole primary ($5.49 \\leq M_{BH} \\leq 8.14~M_{\\odot}$), a high inclination ($60^{\\circ} \\leq i \\leq 71^{\\circ}$), and a large distance ($7.40\\leq d\\leq 12.31$~kpc, Orosz et al. 2001).  Hjellming et al. (2000) monitored the outburst in radio and found jets expanding at apparently superluminal velocities; V4641~Sgr is therefore termed a ``microquasar.'' Significant low-luminosity activity was observed by \\textit{BeppoSAX} and by \\textit{RXTE} Galactic bulge scans (Swank \\& Markwardt 2001) earlier in 1999, however, well before the better-known outburst. In~'t Zand et al. (2000) reported on a pointed \\textit{BeppoSAX} observation made on 1999 March 13.  A line complex was noted, and attributed to Fe~XXV and Fe~XXVI lines with widths narrower than the \\textit{BeppoSAX} resolution.  The best-fit time-averaged 2--10~keV flux measured on March 13 was found to be $2.7\\times 10^{-10}~{\\rm erg}~{\\rm cm}^{-2}~{\\rm s}^{-1}$, corresponding to $1.8\\times 10^{36}~{\\rm erg}~{\\rm s}^{-1} \\leq L_{X} \\leq 4.9\\times 10^{36}~{\\rm erg}~{\\rm s}^{-1}$.  Motivated by the constraints on the accretion flow geometry made through the study of broad lines in other sources at higher luminosities, we re-visit this intermediate-luminosity spectrum of V4641~Sgr. ",
        "conclusions": "We have re-analyzed the \\textit{BeppoSAX} NFI spectra of the Galactic microquasar V4641~Sgr obtained on 1999 March 13.  Applying the Laor model for line emission from the inner part of an accretion disk around a black hole, we find that a broad, relativistically-shaped Fe~K$\\alpha$ line is preferred over a sum of narrow helium-like and hydrogenic Fe~K$\\alpha$ lines at the 4.5$\\sigma$ level of confidence. A line model --- whether a complex of narrow lines or a broadened line --- is strongly required by the data.  Within the context of the Laor model, the narrow emission observed near 6.8~keV (see Figure 1) is the blue wing of a relativistically-shaped line from more neutral Fe species including Fe~I~K$\\alpha$ at 6.40~keV (see Figure 2, Section 3).  The line is among the strongest yet observed in a Galactic black hole system: $W_{K \\alpha} = 870^{+100}_{-70}$~eV.  A spinning black hole (implied if $R_{in} < 6~R_{g}$) is allowed but not required by the line model; our best-fit value is $R_{in} = 8^{+4}_{-3}~R_{g}$. Although the line may be broadened partially by Compton processes, it is likely that relativistic shifts dominate as the profile is strongly asymmetric.  In most cases, observations of Galactic black holes at $L_{X} \\simeq 10^{36}~{\\rm erg}~{\\rm s}^{-1}$ do not achieve the sensitivity required to clearly detect Fe~K$\\alpha$ emission lines or fast quasi-periodic oscillations ($few \\times 100$~Hz; plausibly associated with the Keplerian frequency near the marginally stable orbit).  The Fe~K$\\alpha$ emission line we have observed in V4641~Sgr, then, represents a rare opportunity to test accretion flow models.  Esin, McClintock, \\& Narayan (1997) predict that at $\\dot{m}_{Edd.} \\leq 0.08$, the inner accretion flow in Galactic black holes will be radiatively inefficient and the disk will be recessed to $R_{in} \\geq 200~R_{g}$.  The line we have observed, however, is consistent with being produced in a disk with $R_{in} = 8^{+4}_{-3}~R_{g}$.  In the 0.4--10.0~keV band, we measure $\\dot{m}_{Edd.} = 3.0^{+1.6}_{-1.3} \\times 10^{-2}$ (assuming the most likely values for the black hole mass and distance measured by Orosz et al. 2001).  This value approximately doubles on the 0.4--120~keV range; the value of $\\dot{m}_{Edd.}$ given by a bolometric luminosity might be a factor of a few higher.  We also note that taking extremes for the black hole mass, distance, and observed flux can yield values of $\\dot{m}_{Edd.}$ that are significantly lower than than the range we quote, and also above the critical value of $\\dot{m}_{Edd.} = 0.08$.  It is clear, however, that in the regime where an ADAF is expected to begin to take hold, the observed Fe~K$\\alpha$ line profile strongly suggests a disk which extends close to the marginally stable orbit.  The Fe~K$\\alpha$ line we have measured is not the strongest yet observed from V4641~Sgr.  \\textit{RXTE}/PCA spectra obtained during the decline of the 12.2~Crab flare ($L_{X} \\sim 3-4 \\times 10^{39}~{\\rm erg}~{\\rm s}^{-1}$) in 1999 September, reveal an Fe~K$\\alpha$ line at $\\sim 6.70$~keV with $W_{K \\alpha} \\sim 2.4$~keV (Revnivtsev et al. 2002).  The resolution of the PCA is not well-suited to reveal a line profile of the kind we find in the 1999 March NFI spectra.  In part, the large equivalent width we have measured may be explained by an ionized disk (Ballantyne, Ross, \\& Fabian 2002).  We also speculate that Fe may be over-abundant (relative to solar) by a factor of a few in this source.  Orosz et al. (2001) report that N and Ti may be over-abundant by a factor of 10, Mg by a factor of 7, O by a factor of 3, and Ca by a factor of 2.  The Fe~K$\\alpha$ emission line observed in an \\textit{XMM-Newton} spectrum of the Galactic black hole candidate XTE~J1650$-$500 extends to $\\sim$4~keV and implies a spinning black hole (Miller et al. 2002b).  The line profile seen in the NFI spectra of V4641~Sgr is not as strongly skewed (perhaps extending to $\\sim$5~keV; see Figure 1) as the profile observed in XTE~J1650$-$500, likely indicating that the line is shaped less by strong gravitational effects.  In an NFI observation of the Galactic microquasar GRS~1915$+$105, Martocchia et al. (2002) find a relativistically-shaped Fe~K$\\alpha$ line consistent with neutral ion species, which does not require a spinning black hole (based on the implied inner extent of the accretion disk). Qualitatively, the line profile observed in GRS~1915$+$105 may be the most similar to that in V4641~Sgr.  \\textit{RXTE}/PCA spectra of the Galactic black hole candidate GX~339$-$4 at $L_{X} \\sim few~\\times 10^{36}~{\\rm erg}~{\\rm s}^{-1}$ also reveal evidence for broad, relativistically-shaped line profiles (Nowak, Wilms, \\& Dove 2002) that render an ADAF geometry unlikely."
    },
    "0206/astro-ph0206137_arXiv.txt": {
        "abstract": "{A photometric monitoring campaign of brown dwarfs in the Cha\\,I star forming region in the i and R band revealed significant periodic variations of the three M6.5--M7 type brown dwarf candidates Cha\\,H$\\alpha$\\,2, 3 and 6 (Joergens et al. 2002). These are the first rotational periods for very young (1--5\\,Myr) brown dwarfs and among the first rotational periods for brown dwarfs at all. The relatively long periods of 2.2 to 3.4 days as well as $v \\sin i$ values (Joergens \\& Guenther 2001) indicate that our targets are moderately fast rotators in contrast to very rapidly rotating old brown dwarfs. The periods for the Cha\\,I brown dwarf candidates provide valuable data points in an as yet, in terms of rotational characteristic, almost unexplored region of the age-mass diagram. A comparison with rotational properties of older brown dwarfs indicates that most of the acceleration during the contraction phase takes place within the first 30\\,Myr or less of the lifetime of a brown dwarf. We have also determined periods for the two M5--M5.5 type very low-mass stars B\\,34 and CHXR\\,78C. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206247_arXiv.txt": {
        "abstract": "By comparing the results of numerical microlensing simulations to the observed long-term variability of quasars, strong upper limits on the cosmological density of compact objects in the mass range $10^{-2}$--$10^{-4} \\ M_\\odot$ may be imposed. Using recently developed methods to better approximate the amplification of large sources, we investigate in what way the constraints are affected by assumptions concerning the size of the optical continuum-emitting region of quasars in the currently favored ($\\Omega_\\mathrm{M}=0.3$, $\\Omega_\\Lambda=0.7$) cosmology. ",
        "introduction": "Although the optical variability of quasars on time scales of a few years could be both intrinsic in nature (e.g. due to accretion disk instabilities or supernova explosions) and caused by microlensing along the line of sight, adding the effects of both mechanisms can only increase the probability of variations. By assuming that {\\it all} variability is due to microlensing and comparing the predictions from microlensing scenarios to the observed amplitudes of quasar light curves, upper limits to the cosmological density of compact object may therefore be imposed. This technique was first implemented in \\cite{zackrissoneref1} to constrain compact dark matter populations in the mass range 1--10$^{-4} \\ M_\\odot$ for an Einstein--de Sitter universe. Here, we extend the analysis to the $\\Omega_\\mathrm{M}=0.3$, $\\Omega_\\Lambda=0.7$ cosmology and the full range of quasar sizes allowed. ",
        "conclusions": "By comparing the results of numerical microlensing simulations to the observed long-term variability of quasars, strong upper limits may be placed on the cosmological density of dark matter in the form of compact objects. In the currently favored $\\Lambda$-dominated universe ($\\Omega_\\mathrm{M}=0.3$, $\\Omega_\\Lambda=0.7$), compact objects with masses $\\sim10^{-2}$ $M_\\odot$ and $\\sim10^{-3}$ $M_\\odot$ cannot contribute more to the cosmological density than $\\Omega_\\mathrm{compact} = 0.1$ and $\\Omega_\\mathrm{compact} = 0.05$ respectively,  provided that the continuum-emitting region of quasars is not significantly larger than $1\\cdot10^{13}$ m. We do however note that the method used is sensitive to the assumptions concerning the size of the continuum-emitting region of quasars, and becomes effectively useless for $R_\\mathrm{QSO} \\ge 5\\cdot10^{13}$ m.  The large source amplification formula developed in \\cite{zackrissoneref4} is shown not to have any significant impact in the part of parameter space for which the long-term variability method provides the most interesting constraints.  A more detailed analysis of the uncertainties present in this technique to constrain the cosmological density of compact objects is currently underway (\\cite{zackrissoneref5})."
    },
    "0206/astro-ph0206071_arXiv.txt": {
        "abstract": "We analyze X-ray spectra of heavily obscured ($N_H > 10^{24} \\psc$) active galaxies obtained with \\chandra{}, concentrating on the iron K$\\alpha$ fluorescence line. We measure very large equivalent widths in most cases, up to 5 keV in the most extreme example. The geometry of an obscuring torus of material near the active galactic nucleus (AGN) determines the Fe emission, which we model as a function of torus opening angle, viewing angle, and optical depth. The starburst/AGN composite galaxies in this sample require small opening angles. Starburst/AGN composite galaxies in general therefore present few direct lines of sight to their central engines. These composite galaxies are common, and their large covering fractions and heavy obscuration effectively hide their intrinsically bright X-ray continua. While few distant obscured AGNs have been identified, we propose to exploit their signature large Fe K$\\alpha$ equivalent widths to find more examples in X-ray surveys. ",
        "introduction": "The strongest line in the X-ray spectrum of an active galactic nucleus (AGN) at moderate energies ($4\\lesssim E \\lesssim 10$ keV) is due to iron fluorescence, particularly ``neutral'' Fe K$\\alpha$  at 6.4 keV, from Fe less ionized than \\ion{Fe}{17}. If both the continuum and the Fe-emitting region are viewed directly (the ``Type 1\" view), the equivalent width (EW) is small, typically less than 200 eV. As \\cite{Kro87} pointed out, the EW can increase greatly if the fluorescing material is exposed to a stronger continuum than the observer detects.  In this case (the ``Type 2\" view), an obscuring ``torus'' of material near the active nucleus blocks direct views of the central engine along the line of sight. X-ray observations of Seyfert galaxies generally support this unification scenario. Seyfert 1s typically exhibit EW $\\approx 150$ eV \\citep{Nan94}, while the EWs in Seyfert 2s reported hitherto tend to be larger and distributed more broadly, ranging from about 100 eV to 1 keV \\citep{Tur97}.  Seen from the ``Type 2\" view, the K$\\alpha$ EW depends strongly on the torus geometry and total column density. Previous theoretical calculations \\citep*{Awa91,Ghi94,Kro94} have concentrated on producing EW $\\lesssim 1$ keV, consistent with earlier observations.  They indicated, however, that still larger EWs might result from Compton-thick tori (i.e., $N_H > 10^{24}\\psc$ toward the nucleus) with special geometries.  Motivated by this suggestion and the recent discovery in \\chandra{} X-ray Observatory (\\chandra) observations of several much stronger lines \\citep[e.g.,][]{Sam01}, in this work, we examine the Fe K$\\alpha$ properties of active galaxies that have previously been identified as Compton thick. Our sample, listed in Table \\ref{tab:ew}, comprises all such galaxies \\chandra{} has observed for which published results, archival data, or proprietary data are available. Most of these are classified as Seyfert 2 galaxies. NGC 4945 lacks the requisite optical emission line signature of Seyfert galaxies, but X-rays reveal its active nucleus \\citep{Iwa93}. M51 is exceptional in this group for its low luminosity, and the source identified in the \\chandra{} Deep Field South, CDF-S 202, is exceptional for its high luminosity, but they both fulfill the broad selection criteria. Although this sample is not complete and the data are heterogeneous, these measurements suggest that extremely large Fe equivalent widths are common in such heavily obscured AGN. ",
        "conclusions": ""
    },
    "0206/astro-ph0206500_arXiv.txt": {
        "abstract": "The mean metallicity of the Milky Way thin disc in the solar neighbourhood is still a matter of debate, and has recently been  subject to upward revision (Haywood, 2001). Our star sample was drawn from a set of solar neighbourhood dwarfs with photometric metallicities. In a recent study, Reid (2002) suggests that our metallicity calibration, based on Geneva photometry, is biased. We show here that the effect detected by Reid is not a consequence of our adopted metallicity scale, and we confirm that our findings are robust. On the contrary, the application to Str\\\"omgren photometry  of the Schuster \\& Nissen metallicity scale is problematic. Systematic discrepancies of about 0.1 to 0.3 dex affect the photometric metallicity determination of metal rich stars, on the colour interval 0.22$< b-y <$ 0.59, i.e including F and G stars. For F stars, it is shown that this is a consequence of a mismatch between the standard sequence $m_1(b-y)$ of the Hyades used by Schuster \\& Nissen to calibrate their metallicity scale, and the system of Olsen (1993, 1994ab). It means that although Schuster \\& Nissen calibration and Olsen photometry are intrinsically correct, they are mutually incompatible for metal rich, F-type stars. For G stars, the discrepancy is most probably the continuation of the same problem, albeit worsened by the lack of spectroscopic calibrating stars. A corrected calibration is proposed which renders the calibration of Schuster \\& Nissen applicable to the catalogues of Olsen. We also give a simpler calibration referenced to the Hyades sequence, valid over the same colour and metallicity ranges. ",
        "introduction": "Solar neighbourhood stars serve as reference to which we compare the characteristics of the Galaxy outside the immediate solar vicinity, and their properties scale our measurements of the galactic  structure and evolution. In view of their importance, it is somehow surprising that their general properties, such as the mean metallicity of the galactic disc stars, is still a matter of debate. In a recent paper (Haywood 2001), we constructed a metallicity distribution from stars within 20pc from the sun. This new metallicity distribution was shown to be centred on solar metallicity, or 0.1-0.2 dex higher than the value found in most previous studies. We discussed that this discrepancy is the result of various biases that enter the definition of these samples, the principal effect being caused by the selection of samples on the basis of spectral type. Another sensitive effect comes from the adopted metallicity scale. The choice of a given metallicity scale is determined from two criteria : the necessity to utilise  stars with as low masses as possible, in order to avoid biases favouring young stars, and available photometry. The most widely  used photometry  for studying the metallicity distribution is the Str\\\"omgren photometry, with the metallicity scale from Schuster \\& Nissen (1989) (hereafter SN). In our own study, and  while the calibration of SN is given as valid down to $b-y$=0.59 (B-V$\\approx$1.0), we used Geneva photometry and  the metallicity scale from Grenon (1978) for stars redder than $b-y\\approx0.42$ (B-V$\\approx$0.67). The reason for this choice was that Geneva photometry is available for a larger set of solar neighbourhood K dwarfs. Hence, in the initial sample used by Haywood (2001), approximately half the stars had their metallicity determined from Str\\\"omgren photometry, the other half from Geneva photometry.   In a recent paper, Reid (2002) finds that the (B-V,[Fe/H]) distribution of our sample shows a trend of about 0.2~dex from B-V=0.5 to B-V=1.0, suggesting  that the calibration from Grenon (1978) is plagued by a systematic error. By constrast, his  (B-V, [Fe/H]) distribution, based entirely on the metallicity scale of SN, shows no such trend, if characterized by a simple linear regression. Most recently, and while terminating our paper, a study on the same subject was presented by Twarog, Twarog \\& Tanner (2002), exactly pointing to the problem we had discovered in Reid (2002) and that motivated the present work - a severe apparent deficiency in the metallicity scale of SN, causing underestimates of the photometric metallicity of late G type and early K type metal rich stars.   It appears however that the study by Twarog et al. has underestimated the discrepancy between photometric and spectroscopic abundances, in the sense that this problem also affects F-type stars, in contrast with their claim. Moreover, while they suggest that the origin of this discrepancy is an unsuspected high dependence of the metallicity on the $c_1$ index, we propose that it probably originates in a mismatch between the $m_1(b-y)$ standard sequence adopted in SN - which comes from Crawford (1975) - and the one that can be deduced from the photometric catalogues of Olsen (1993, 1994ab) (see also Olsen (1984), table VI), which constitute the vast majority of available measurements of Str\\\"omgren  photometry. The aim of the present paper is threefold :  (1) In section 2, we first present the discrepancy and quantify its amplitude on the whole colour range, that is between  0.22$< b-y <$ 0.59.  (2) In section 3, we propose an explanation for the origin of this effect, and give a corrected calibration, for which we calculate new coefficients on the basis of an enlarged spectroscopic dataset. Using the same spectroscopic calibrating stars, we also give an alternative calibration using $\\delta m_1$ and $\\delta c_1$.  (3) Section 4 is a brief discussion of the impact of  the photometric calibration on the metallicity distribution of the solar neighbourhood stars. In particular, it is demonstrated that the part of the sample in Haywood (2001) that is concerned by this effect is unchanged by the new calibration. That means that the results presented in Haywood (2001) are robust.    \\begin{figure*} \\includegraphics[width=16cm]{fig1.ps} \\caption{Metallicity distribution as a function of B-V for our sample (a), and Reid (2002) sample (b). Our sample is a mixture of stars with Str\\\"omgren metallicity (squares) and Geneva metallicity (circles). The horizontal line represents the metallicity of the Hyades cluster from Perryman et al. (1998), at [Fe/H]=+0.14. The star symbols  are Hyades members from de Bruijne et al. (2001), with metallicity determined from the calibration of Grenon (1978). The region delimited on the right lower part of the plot is the region where Favata Micela \\& Sciortino (1997) found no objects in the their sample. In the sample of Reid (2002) (plot b), all stars have their metallicity determined from Str\\\"omgren photometry, through Schuster \\& Nissen (1989) calibration. Star symbols also represent Hyades members for which metallicity has been calculated through the calibration of  Schuster \\& Nissen (1989). Our Str\\\"omgren metallicities in plot (a) were determined with the calibrations of SN but corrected as described in Haywood (2001), and are not strictly equal to those of Reid (2002), although patterns common to the two ($B-V$,$[Fe/H]$) distributions can be seen. } \\end{figure*} ",
        "conclusions": "Using the new calibration, it is now possible to recalculate  metallicities for the set of local dwarfs. The new (B-V, [Fe/H]) distribution is shown on Fig. 10. The result is now satisfactory, but the occasion is taken to emphasise the need for a larger number of spectroscopic metallicities for cool dwarfs.  In the last decade, the calibration of SN have been widely used in various studies, in particular to design metallicity distribution of long-lived dwarfs in the solar neighbourhood. In one such study, we  sampled the solar neighbourhood within 20~pc (Haywood, 2001). The sample, before the selection of long-lived dwarfs, contained 177 stars for which metallicity came from Geneva photometry, and 172 stars for which metallicity was calculated from Str\\\"omgren metallicity, mostly for stars bluer than B-V=0.67.  Str\\\"omgren photometry was also used for 41 stars redder than this limit, for which no Geneva photometry was available. All stars with SN metallicities were corrected as  [Fe/H]$_{SN}$/0.865+0.052. Only 20 stars with colour in the critical interval 0.8$<B-V<$1.0 were included in the sample. In view of the limited number of objects affected by the problem, it is unlikely that our metallicity distribution has suffered much from this effect. We conclude that the trend seen on  Fig. 1a between  metallicity  and colour is obviously real.  As a last check on our metallicities, we have searched for spectroscopic iron abundances for the stars in Fig. 1a in the catalogue of Cayrel et al. (2001). Fig. 11  shows both the photometric and spectroscopic values $vs$ B-V colour and there is  good general agreement between the two. The spectroscopic metallicities, although still very sparse, confirm the general shape of the B-V-[Fe/H] distribution, and we conclude that our findings of Haywood (2001) are robust.   \\begin{figure*} \\includegraphics[width=16cm]{fig11.ps} \\caption{B-V-[Fe/H] distribution for the sample in Haywood (2001).  Circles represent stars for which spectroscopic measurements are available from the catalogue of Cayrel et al. (2001). The photometric measurements are shown as squares, and related to the spectroscopic value with a vertical line. This plot shows that there is no systematic deviation between spectroscopic metallicities and our photometric abundances. } \\end{figure*}"
    },
    "0206/astro-ph0206446_arXiv.txt": {
        "abstract": "K-dwarfs have life-times older than the present age of the Galactic disc, and are thus ideal stars to investigate the disc's chemical evolution.  We have developed several photometric metallicity indicators for K dwarfs, based an a sample of accurate spectroscopic metallicities for 34 disc and halo G and K dwarfs. The photometric metallicities lead us to develop a metallicity index for K dwarfs based only on their position in the colour absolute-magnitude diagram. Metallicities have been determined for 431 single K dwarfs drawn from the Hipparcos catalog, selecting the stars by absolute magnitude and removing multiple systems. The sample is essentially a complete reckoning of the metal content in nearby K dwarfs. We use stellar isochrones to mark the stars by mass, and select a subset of 220 of the stars which is complete in a narrow mass interval. We fit the data with a model of the chemical evolution of the Solar cylinder. We find that only a modest cosmic scatter is required to fit our age metallicity relation. The model assumes two main infall episodes for the formation of the halo-thick disc and thin disc respectively. The new data confirms that the solar neighbourhood formed on a long timescale of order 7 Gyr. ",
        "introduction": "A central issue in studies of the chemical evolution of the Galactic disc is to resolve the so called ``G-dwarf problem'' (van den Bergh, 1962; Schmidt, 1963; Pagel and Patchett, 1975). The problem is that the observed stellar metallicity distribution shows far fewer metal deficient stars than the predictions of the simplest, closed box models of the Galactic disc's chemical evolution. Integrated light studies indicate that the G-dwarf problem is not restricted to our own disc (Worthey et al, 1996) but is also found in other galaxies. There are many ways that the evolutionary models can be modified to bring them into consistency with the observations, such as pre-enrichment of the gas, a time dependent Initial Mass Function or gas infall (for a review see e.g. Pagel, 1997)  G-dwarfs are sufficiently massive that some of them have begun to evolve away from the main sequence, and these evolutionary corrections must be taken into account when determining their space densities and metallicities. While these problems are by no means intractable, it has been long recognized that K dwarfs would make for a cleaner sample of the local metal abundance distribution, because for these stars the evolutionary corrections are negligible.  K dwarfs are of course intrinsically fainter, and it has not been until recently that accurate spectroscopic K dwarf abundance analyses have become available, with which to calibrate photometric abundance indicators (Flynn and Morell, 1997). As a result, it is now known that there is a K-dwarf problem which is very similar to the G-dwarf problem (Flynn and Morell, 1997).  Studies of M dwarfs indicate that the problem is present in these stars too (Mould 1976). Such stars are still not really ideal for measuring the local metallicity distribution, because metallicities for M dwarfs are difficult to determine in the optical (but appears possible using infrared $JHK$ photometry, (see e.g. Stauffer and Hartmann, 1986; Leggett et al, 2000). The all sky surveys presently underway in the infrared (Denis, 2MASS) may make such stars viable metallicity tracers in the near future.  The G-dwarf metallicity distribution is already an extensively studied subject (Pagel and Patchett, 1975; Sommer-Larsen, 1991; Wyse and Gilmore, 1995; Rocha-Pinto and Maciel, 1996) and has been used by chemical evolution models to constrain the time scale of the formation of the Galactic disc at the solar neighbourhood. The G-dwarf metallicity distribution can be well fit by models in which gas has been settling onto the disc over a protracted period, of some billion years.  For example, Chiappini et al (1997), on the basis of the fit of the G-dwarf metallicity distribution, have shown that the disc in the solar neighbourhood was formed on a long time scale of 7 - 8 Gyr. This conclusion has been later stressed also by other authors (e.g. Portinari et al, 1998; Prantzos and Silk, 1998; Chang et al, 1999).  The release of the data from the European Space Agency's {\\em Hipparcos} (ESA, 1997) satellite offers a great opportunity to determine the local metallicity distribution of the disc from a complete sample of K dwarfs. There are three clear improvements because of Hipparcos. Firstly, the parallaxes are so accurate that the K dwarfs can be selected by absolute magnitude rather than colour, which is a much better way of isolating stars in a particular mass range. Secondly, the uniformity of the Hipparcos data allows us to construct samples with well understood completeness limits. Thirdly, a large fraction of the close binaries can be identified using Hipparcos and removed (since photometric abundance indicators are calibrated for single stars). This last effect turns out to be quite important.  In this paper we develop a number of simple photometric metallicity indicators for K dwarfs, based on a spectroscopically determined sample of metallicities by Flynn and Morell (1997). The metallicities for several hundred K dwarfs drawn from the Hipparcos catalog have been measured via new observations described here. This sample has been used to show that there is a very simple relation between the $V$ band luminosity $M_V$ of K dwarfs at a given $B-V$ colour and metallicity (the work is fully described in a companion paper (Kotoneva, Flynn and Jimenez, 2002, paper II). This simple relation is used in this paper to determine metallicities for 431 single K dwarfs in a near complete sample selected from the Hipparcos catalog. We examine a set of isochrones and find a simple relation between the position of a star in the Hipparcos colour magnitude diagram and mass. We are thus able to select a sample of dwarf stars by the more physically relevant parameter of mass, rather than by colour or absolute magnitude, as is usually the case. The metallicity distribution obtained is then compared to the results of models of the chemical evolution of the Solar neighbourhood.  The paper is organized as follows: In section 2 we describe the selection of the K dwarf sample and in section 3 the observations and reductions of the stars in broadband and intermediate band photometric systems. The details of the various metallicity calibrations are presented in section 4. In section 5 we use isochrones to fit the masses of the stars, and select the stars in an appropriate mass range which is near complete; we discuss the kinematics of the stars and compute corrections for the metallicity distribution function from the solar volume to the solar cylinder.  The metallicity distribution is compared briefly with other determinations in the literature. We compare the data to the predictions of a model of the chemical evolution of the Solar neighbourhood due to Chiappini et al (2001) in section 6. We draw our conclusions and summarize in section 7. ",
        "conclusions": "We have calibrated several photometric abundance indices for K dwarfs based on a sample of 34 G and K dwarfs with accurate, spectroscopically determined metallicities. Two of the indices use Cousins $R -I$ photometry to estimate stellar effective temperature and the Geneva $b_1$ or Str{\\\"o}mgren $m_1$ and $c_1$ colours.  These indices give metallicity estimates of $\\approx 0.2$ dex accuracy. A third metallicity index uses $B-V$ and $V-I$ broadband colours. A fourth metallicity index, described in detail in a companion paper (Kotoneva, Flynn and Jimenez 2002, paper II), is based on absolute magnitude $M_V$ and $B-V$ colour. These latter two indices appear to provide metallicities of $\\approx 0.1$ dex accuracy.  For a set of newly acquired observations of K dwarfs, we use one of these indices to obtain the metallicity distribution of a near complete sample of K dwarfs in the Solar neighbourhood drawn from the Hipparcos catalog. Care has been taken to remove the multiple stars from the sample, for which metallicities cannot be measured accurately. Through isochrones we assign masses to the K dwarfs, and select those K dwarfs which fall into a mass window $0.75 < $ M/M$_\\odot < 0.83$, within which the sample is near to complete. This yields a sample of 220 K dwarfs.  The metallicity distribution of the 220 K dwarfs is strongly peaked near the solar metallicity, confirming the existence of the ``G dwarf problem'' amongst K dwarfs, as seen in several earlier studies (Marsakov and Shevelev (1988), Flynn and Morell (1997) and Rocha-Pinto and Maciel (1998).  We compare the metallicity distribution with Galactic chemical evolution models of Chiappini et al (2001). In the context of this model of the metallicity evolution, we find that the amount of cosmic scatter in the metallicities is small, not more than 0.15 dex. The model match the data well, indicating that the disc was formed via infall processes over an extended time scale of order 7 Gyr."
    },
    "0206/astro-ph0206393_arXiv.txt": {
        "abstract": "We present a model for star formation and supernova feedback that describes the multi-phase structure of star forming gas on scales that are typically not resolved in cosmological simulations.  Our approach includes radiative heating and cooling, the growth of cold clouds embedded in an ambient hot medium, star formation in these clouds, feedback from supernovae in the form of thermal heating and cloud evaporation, galactic winds and outflows, and metal enrichment. Implemented using smoothed particle hydrodynamics (SPH), our scheme is a significantly modified and extended version of the grid-based method of Yepes et al. (1997), and enables us to achieve high dynamic range in simulations of structure formation.  We discuss properties of the feedback model in detail and show that it predicts a self-regulated, quiescent mode of star formation, which, in particular, stabilises the star forming gaseous layers of disk galaxies.  The parameterisation of this mode can be reduced to a single free quantity which determines the overall timescale for star formation.  We fix this parameter numerically to match the observed rates of star formation in local disk galaxies.  When normalised in this manner, cosmological simulations employing our model nevertheless overproduce the observed cosmic abundance of stellar material.  We are thus motivated to extend our feedback model to include galactic winds associated with star formation.  Using small-scale simulations of individual star-forming disk galaxies, we show that these winds produce either galactic fountains or outflows, depending on the depth of the gravitational potential. In low-mass haloes, winds can greatly suppress the overall efficiency of star formation.  When incorporated into cosmological simulations, our combined model for star formation and winds predicts a cosmic star formation density that is consistent with observations, provided that the winds are sufficiently energetic.  Moreover, outflows from galaxies in these simulations drive chemical enrichment of the intergalactic medium, in principle accounting for the presence of metals in the Lyman alpha forest. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\footnotetext[1]{E-mail: volker@mpa-garching.mpg.de} \\footnotetext[3]{\\hspace{0.03cm}E-mail: lars@cfa.harvard.edu}%  The past two decades have witnessed the emergence of a successful class of theoretical models to explain the formation of structure in the Universe.  In these ``hierarchical'' scenarios, cold dark matter (CDM) forms the dominant mass component, and structure grows via gravitational instability from primordial perturbations seeded by inflation.  Galaxy formation proceeds ``bottom-up'' in this picture, with smaller galaxies forming first, and larger objects being assembled through merging of less massive building blocks.  Numerical simulations have become an important tool for exploring the detailed predictions of these models.  For example, the clustering properties of matter on large scales, where gasdynamical forces are negligible compared to gravity, can now be computed with high precision.  However, extending these studies to understand the formation of luminous components of galaxies is considerably more difficult.  Cosmological simulations that attempt to include the essential physical processes associated with galaxy formation, namely radiative cooling and star formation, invariably encounter formidable difficulties.  For this reason, the most popular approach to date for describing hierarchical galaxy formation has been the so-called ``semi-analytic'' technique \\citep[e.g.][]{Wh91,Lac91,Kau93a,Kau94,Col94}, in which highly simplified prescriptions for the baryonic physics are combined with Monte-Carlo methods to construct merging history trees.  In recent times, these Monte-Carlo trees have been replaced with merging histories obtained directly from high-resolution N-body simulations \\citep{Kau99,Kau99b,Spr99b}.  Some impressive successes in direct numerical simulation of cosmological hydrodynamics have been obtained since this method was pioneered more than a decade ago \\citep{Ev88,Ev90,He89b,He89,Nav91,Ba91,Ka91,Hi91,Th92,KaHeWe92,Ce93}. For example, these studies improved our understanding of quasar absorption line systems \\citep[e.g.][]{CMOR94,Zh95,He96}, the intergalactic medium \\citep[e.g.][]{Ce93,Ce00,Ka96,Kat99,We97,Wein2000,St95,Bry94,Bry98,Bl99, Pea99,Pea2000}, and galaxy formation \\citep[e.g.][]{Ka91,Nav94,St95,Mi96, Wa96,Na97,St99}.  However, the numerical modelling of key processes related to star formation and associated feedback by supernovae and stellar winds has remained problematic.  Because radiative cooling is particularly efficient in dense, low-mass haloes at high redshift, simulations are confronted with an overcooling problem, which is typically manifested through an overproduction of stars.  Semi-analytic treatments indicate that a physical handling of feedback processes is required to regulate star formation and reduce the fraction of baryons that collapse in the centres of haloes.  Unfortunately, it is not clear how best to incorporate feedback processes into simulations of galaxy formation, although they appear to be crucially important to a successful description of hierarchical CDM universes.  Here, we present a model for star formation and feedback that goes beyond a purely phenomenological description.  In most attempts to simulate galaxy formation, the gas within galaxies is treated as a single phase medium and is converted into stars on some characteristic timescale, which is usually related to the local dynamical time. Often, auxiliary criteria requiring that the flow be locally converging or Jeans unstable are also invoked to restrict the eligibility of gas fluid elements for star formation.  Based on an assumed stellar initial mass function (IMF), it is then straightforward to work out the expected rate of energy input due to associated supernovae.  However, it is not clear how this ``feedback energy'' should be deposited into the gas.  Simply augmenting the thermal reservoir of the gas has little effect, because the cooling time in the dense single-phase medium in these simulations is so short that the energy is promptly radiated away.  Consequently, pure thermal feedback in this manner fails to solve the overcooling problem.  Other approaches have been attempted to prevent this rapid loss of feedback energy, thereby providing a back-reaction on star formation \\citep[see][for a recent comparison of the most common of these methods]{Kay01}.  For example, \\citet{Th2000} temporarily lower the densities of SPH particles receiving feedback energy to prevent immediate energy loss.  Alternatively, such particles have been treated adiabatically for a brief period of time after acquiring feedback energy \\citep{Ge97}.  Although crude, these methods show some success in regulating star formation.  For example, \\citet{Th2001} have recently found that the angular momentum problem in the formation of disk galaxies can be alleviated through the combined effects of a smoothed deposition of feedback energy with a delay of cooling by 30 Myr.  Another method for depositing the energy is through kinetic feedback, where radial momentum kicks are imparted to particles surrounding sites of star formation \\citep[e.g.][]{Na93,Mi94}, but the fraction of feedback energy that must be put into this channel depends on numerical details of the simulations.  Also, the size of the affected regions depends on spatial resolution, and is in general much larger than the physical scale of individual supernova remnants.  \\citet{SomLar99} experimented with yet more extreme versions of feedback in order to investigate the effects of galactic outflows. For example, in their ``blow-out'' simulations, they identified all dense, cold gas at $z=2.4$ and instantly removed it and redistributed it uniformly in spherical shells around the centres of the gas clumps. While this is clearly a highly artificial model, it demonstrated that outflows can have a large impact on galaxy formation.  In particular, the loss of angular momentum by the gaseous component was greatly reduced in their simulations, alleviating the apparent discrepancy between observed and simulated values of the specific angular momentum of galactic disks. \\citet{Scan01c} used a similar gas-removal technique to model galactic winds at high redshift.  A disadvantage common to many of these recipes is their explicit or implicit dependence on numerical details of the simulations.  For example, if a criterion for star formation explicitly involves the mass of a particle, it is unclear in what sense numerical convergence for the star formation rate of a specific object can be expected if the simulation is repeated with a different mass resolution.  Instead, the free parameters describing star formation may need to be recalibrated whenever the properties of the simulation are altered. \\cite{Spr99} tried to overcome this problem by developing a feedback method based on an effective model for the star-forming interstellar medium (ISM), where turbulent motions on unresolved scales were postulated in an ad-hoc fashion to provide pressure support for the gas, thereby regulating star formation.  In essence, feedback energy was only slowly thermalised in this method, thereby delaying the dissipation of feedback energy in a similar way to some of the models noted above.  Because the Lagrangian nature of smoothed particle hydrodynamics makes it relatively easy to obtain the high spatial dynamic range needed for simulations of galaxy formation, it has so far been the most popular numerical technique for studying this problem.  However, star formation and feedback have also been incorporated into hydrodynamical mesh codes \\citep[e.g.][]{Ce93,Gne98}, but their limited spatial resolution is a serious concern, because star formation occurs at substantially higher overdensities than can be resolved with a fixed grid in cosmological volumes.  This is quite different when high-resolution Eulerian codes are applied to simulations of isolated galaxies, as in the work of \\citet{Mac99}, who carried out detailed studies of starbursts in dwarf galaxies.  Eulerian codes have also been employed to directly resolve the multi-phase structure of the ISM, yet state-of-the-art simulations are either restricted to two dimensions \\citep{Wa99}, or can presently only resolve regions of size $\\sim$100 pc or so in three dimensions \\citep{Wada01}.  \\citet{Ye97} adopted an intermediate approach by developing a sub-grid model of a two-phase ISM and implementing it in an Eulerian code. Their spatial resolution was consequently relatively poor, despite restricting themselves to rather small cosmological volumes, but together with subsequent work \\citep{Eli99b,Eli99a,Asc02} they were able to demonstrate a number of interesting properties of this scheme, which is physically better motivated than most of the models currently employed in cosmological SPH simulations.  Because of this, we will base our analysis in part on their approach and on a first implementation of it in SPH by \\citet{Hu99}, yet we will significantly extend and modify the model.  Our specific aim is to formulate a description of feedback which is physically motivated, numerically well specified, and suitable for simulating galaxy formation within large cosmological volumes.  In what follows, we show that our model leads to self-regulation of star formation, and describe how it can be normalised by analysing its properties.  Under the assumption that the normalisation obtained from observations of disk galaxies in the present Universe holds at all redshifts, we find that the resulting cosmic star history leads to an overprediction of the luminosity density of the universe.  We then argue that galactic winds are a plausible mechanism for eliminating this discrepancy, and they simultaneously help to account for other observational data, like the metals found in the low-density intergalactic medium (IGM), and describe a phenomenological strategy for including these winds in cosmological simulations.  This paper is organised as follows. In Section 2, we summarise our new hybrid multi-phase model for quiescent star formation in detail, and analyse some of its basic properties.  In Section 3, we discuss physical and observational constraints on the values of the model parameters.  In Section 4, we extend the model by including galactic winds.  We then describe our numerical implementation of the method in a parallel TreeSPH code for cosmic structure formation in Section 5. We present numerical results for the formation of disk galaxies in isolation, and for small simulations of cosmic structure formation in Section~6.  Finally, we give conclusions in Section~7. ",
        "conclusions": "We have presented a new model for the treatment of star formation and feedback in cosmological simulations of galaxy formation.  Our approach leads to the establishment of a tight self-regulation cycle for star formation which is based on a rough, yet physically motivated model of the ISM. A crucial ingredient of the model is the assumed multi-phase structure of the ISM. At high density, much of the interstellar material becomes bound in cold clouds, becoming sites for star formation. The density of the ambient phase is lowered accordingly, which also reduces its radiative losses and allows it to be heated by supernovae to a temperature high enough to provide some pressure support for the ISM. The stabilising effect of this pressure counteracts star formation and, together with the cloud formation/evaporation cycle, leads to self-regulation.  Due to the nature of the approximations made, it is clear that our model is still phenomenological to a large degree.  In principle, however, the physical approximations we used can be refined in the future, thereby improving the faithfulness of the model.  The star formation timescale in the quiescent mode of star formation can be directly determined from observations of local disk galaxies. If this normalisation is adopted, we find that cosmological simulations nevertheless lead to an overproduction of stars.  We have invoked galactic winds as a heuristic extension of our model to remedy this problem. While there is mounting observational evidence for the existence of such winds, they are introduced in our model in a phenomenological manner. However, winds clearly have a range of highly interesting consequences. They reduce the global efficiency of star formation, and this suppression is particularly strong in low-mass haloes, thereby helping to explain the observed ``low'' values of the luminosity density. Winds are also capable of accounting for the enrichment of the low-density IGM with metals, and they may be crucial to understanding the metal distribution in the intra-cluster gas of clusters of galaxies.  Another important property of our technique is that it is numerically well posed in the sense that the parameters of the model can be determined directly based on physical arguments or observational input, and need not to be changed when the mass resolution is varied. The use of the sub-resolution technique introduces a well-specified continuum limit in the hydrodynamic equations, despite the inclusion of cooling.  Unlike in standard single-phase simulations, an unphysical collapse of gas to a scale set by the resolution limit of the simulation is thus prevented.  Combined with improved formulations of SPH \\citep{SprHe01} this makes it much easier to reach numerical convergence even under conditions of only moderate resolution.  It will thus be interesting to study the predictions of our model in more detail.  As first steps, we have already carried out a detailed analysis of the predictions of the model for the cosmic star formation history in the $\\Lambda$CDM cosmology \\citep{SprHerSFR}, and we analysed the impact of cooling, star formation, and winds on secondary anisotropies of the cosmic microwave background \\citep{WhiHerSpr02}. In these studies, a new comprehensive set of simulations that included strong galactic winds was used.  We plan to extend our analysis of these simulations in future work."
    },
    "0206/astro-ph0206116_arXiv.txt": {
        "abstract": "{The recent report of iron line features in the afterglow of the gamma-ray burst (GRB) 991216 has important implications for the properties of the radiating material and hence the nature of the immediate burst environment. We argue that the putative strong Fe emission line can be attributed to the reflected emission from Thomson-thick matter which is illuminated by a power-law continuum. The ionization parameter of the material (i.e., the flux to density ratio) is around $10^{3}$, resulting in a Fe K$\\alpha$ line from He-like iron. A supersolar abundance of iron is not required by the data. Interestingly, the ionizing continuum must be harder than the observed one.  We interpret this as due to the fact that the observed continuum is dominated by the standard blast wave emission, while the line is produced by reprocessing material located at much smaller radii. ",
        "introduction": "\\label{sec:int}  The detection of iron emission features in the X-ray spectra of several $\\gamma$-ray burst (GRB) afterglows (GRB~970508: Piro et al. 1999; GRB~970228: Yoshida et al. 2001; GRB~991216: Piro et al. 2000; GRB~000214: Antonelli et al. 2000) is of great importance to the understanding of the nature of the burst emission and particularly the burst progenitor (e.g., M{\\'e}sz{\\'a}ros \\& Rees 1998; Lazzati et al. 1999; B\\\"ottcher \\& Fryer 2001).  All the lines have been detected roughly one day after the burst explosion due to instrumental limitations (the line may have been there at earlier times, but the afterglow was not yet observed), with an equivalent width (EW) of the order of 1~keV and a luminosity $L_{\\rm{Fe}}\\approx10^{44}$~erg~s$^{-1}$.   The large EW inferred from the X-ray features favour models in which the line is produced by reflection from a highly ionized surface, rather than transmission. Several different emission mechanisms and geometries have been proposed to account for the line emission, all of which fall mainly into one of two categories: {\\it geometry} dominated (GD) models or {\\it engine} dominated (ED) models. Both classes associate the scattering medium with debris from a stellar collapse. In the GD models, the X-rays from the burst and early afterglow emission illuminate a mass $M_{\\rm{Fe}}>0.06M_{\\odot}$ of iron material at a distance of about $10^{16}$~cm from the burst location (Lazzati et al. 1999; Vietri et al. 2001; see also Lazzati et al. 2001). This material may be associated with the remnant of a supernova that occurred days to months before the burst (Vietri \\& Stella 1998). Alternatively, in the ED models, the line emission can be attributed to the interaction of a long lasting outflow from the central engine with the progenitor stellar envelope at distances $R \\le 10^{13}$~cm (Rees \\& M{\\'e}sz{\\'a}ros 2000; M{\\'e}sz{\\'a}ros \\& Rees 2001; B\\\"ottcher \\& Fryer 2001). In this case, only a small mass of Fe is required and there is no need for a pre-ejected supernova shell (i.e., the line can be explained if the burst and the star explode simultaneously; see Woosley 1993).  There are several ways in which these two models can be distinguished. First, the line produced by a compact reprocessor will show a higher degree of variability. It is likely that in any ED scenario the line intensity will decrease with time, mirroring the decay of energy input. Due to geometrical time dilution of the line photons, it is hard for any GD model to produce line variability on a time scale shorter than $t_{\\rm var}<R/c\\approx1$~d. Unfortunately, the detected lines were too weak to allow for a proper variability analysis, and no firm conclusion could be made (while a hint of non-variability was derived in GRB~000214 [Antonelli et al. 2000], the line in GRB~970508 seemed to disappear before the end of the observation [Piro et al. 1999]). Secondly, different ionization parameters $\\xi\\equiv L_X/(n_H\\,R^2)$ are expected, with $\\xi\\approx 10^3$ for ED models and $\\xi\\approx 10^5$ for GD models. This would suggest that ED models may produce a more luminous line (see Ballantyne \\& Ramirez-Ruiz 2001), but smaller ionization parameters can be predicted even for the GD models if, for example, the emitting material is clumped or the off-axis burst emission is less intense than the one we receive. The equivalent widths of the lines may also be very useful. In the ED models it is impossible to avoid seeing the ionizing continuum together with the reflected component, while it is natural in GD models to see only the reflected component. This implies (Ballantyne, Fabian \\& Ross 2002) that a line with EW$>800$~eV can be produced only by GD models irrespective of the iron abundance.  Finally, broad-band X-ray spectroscopy could help settle this question. This is because in the ED case, the line should be due to material newly synthesized by the exploding star (which is mostly nickel rather than iron; see Woosley \\& Weaver 1995), while in the GD case, the star exploded several months before the GRB and the nickel has had time to decay into iron. It is, however, possible for nickel to be bypassed due to high neutronization (where the iron is directly synthesized by the exploding star), or by the iron being dragged out from the stellar core by the expanding fireball.  Measurements of the strength and shape of the reflection spectrum can yield the geometry, velocity, density and abundances of the scattering medium. In this paper, we attempt to fit detailed models of such reflection spectra, computed with the method described by Ross \\& Fabian (1993), to the X-ray afterglow of GRB~991216. This will allow, for the first time, a self-consistent determination of both the Fe~K$\\alpha$ line and the continuum in a X-ray afterglow of a GRB. A description of the data and the fitting results are presented in \\S 2. In \\S 3 we briefly discuss the general constraints that arose from the fits in the context of both GD and ED models. ",
        "conclusions": "\\begin{tabular}{@{}llllllllll} Fit \\# & Model & $\\Gamma_1$ & $N_{\\mathrm{H}}$ & $E$ & $\\sigma$ & $\\Gamma_2$ & $\\chi^2/$d.o.f. & Sig. & Notes \\\\ \\hline 1 & PL & 1.76$^{+0.09}_{-0.10}$ & -- & -- & -- & -- & 99/60 & -- & -- \\\\ 2 & A*PL & 1.83$^{+0.14}_{-0.13}$ & 0.24$^{+0.40}_{-0.24}$ & -- & -- & -- & 97/59 & 74\\% & -- \\\\ 3$^{\\dag}$ & PL & 1.83$\\pm 0.11$ & -- & -- & -- & -- & 85/50 & -- & -- \\\\ 4$^{\\dag}$ & A*PL & 2.03$^{+0.11}_{-0.19}$ & 0.52$^{+0.55}_{-0.40}$ & -- & -- & -- & 80/49 & 92\\% & -- \\\\ 5$^{\\dag}$ & PL+G & 1.87$\\pm 0.11$ & -- & 6.80$^{+0.02}_{-0.19}$ & 0.01$^{\\mathrm{fixed}}$ & -- & 72/48 & 98\\% &  -- \\\\ 6$^{\\dag}$ & A*(PL+G) & 2.13$^{+0.23}_{-0.21}$ & 0.63$^{+0.58}_{-0.43}$ & 6.80$^{+0.02}_{-0.19}$ & 0.01$^{\\mathrm{fixed}}$ & -- & 65/47 & 97\\% & EW $\\approx$ 300~eV \\\\ 7$^{\\dag}$ & A*(PL+G) & 2.17$^{+0.26}_{-0.21}$ & 0.68$^{+0.63}_{-0.44}$ & 6.82$^{+0.10}_{-0.11}$ & 0.15$^{+0.15}_{-0.11}$ & -- & 62/46 & 84\\% & EW $\\approx$ 400~eV \\\\ 8 & same as 6 & -- & -- & -- & -- & -- & 96/57 & -- & -- \\\\ 9 & A*(Bkn. PL+G) & 2.11$^{+0.27}_{-0.19}$ & 0.59$^{+0.63}_{-0.40}$ & 6.80$^{+0.02}_{-0.19}$ & 0.01$^{\\mathrm{fixed}}$ & 0.40$^{+0.70}_{-0.84}$ & 72/55 & -- & $E_{\\mathrm{brk}}=4.2\\pm 0.7$~keV \\\\ 10 & A*(PL+R) & 2.61$^{+0.59}_{-0.40}$ & 1.37$^{+0.84}_{-0.51}$ & -- & -- & 0.0$^{+0.44}_{-0.0p}$ & 76/56 & -- & $\\log \\xi=3.01^{+0.07}_{-0.11}$ \\\\ \\hline \\end{tabular}  \\medskip $_p$ Parameter pegged at lower-limit. $^{\\dag}$ Ignored $E>4.6$~\\kev.  \\end{minipage} \\end{center} \\end{table*} A power-law fit to the entire energy range (Fit \\#1) does not result in an adequate fit. Introducing intrinsic absorption (the Galactic column of $2.1\\times 10^{21}$~cm$^{-2}$ is included in all fits) does not improve this result. The poor fit is a result of a line feature around 3.4~\\kev\\ (observed frame) and of a significantly hardening of the data above $\\sim 5$~\\kev. This was noticed by Piro \\etal\\ (2000) who interpreted the excess above 5~\\kev\\ as an iron recombination edge in emission.  Ignoring this change of slope for the time being, we concentrate on fitting the data below 4.6~\\kev\\ (the fits marked with a $\\dag$). Fitting a power-law to this limited energy range still does not provide a good fit to the data (\\#3). Adding intrinsic absorption to the model (\\#4) only marginally helps the fit, with the added component only significant at the 92\\% level, according to the F-test. Replacing the absorber with a Gaussian emission line (Fit \\#5) does improve $\\chi^2$ substantially. The line is significant at the 98\\% level, and its centroid energy is tightly constrained around the energy of the He-like \\fe\\ line\\footnote{However, if we ignore the change in continuum slope and fit a power-law plus Gaussian line model to the entire energy range, the upper-limit of the line energy is 6.98~keV at $z=1.02$ (roughly the error on the measured redshift). Therefore, we cannot strictly rule out a contribution from H-like iron.}.  Including intrinsic absorption in this model (Fit \\#6) improves the fit still further. Allowing the width of the Gaussian to vary (Fit \\#7) does not significantly improve the fit, but does show that the line is marginally resolved and is likely broadened. Therefore, the true error on the centroid energy is $\\pm 0.1$~\\kev, as the narrow line fit above missed a second minimum in $\\chi^2$ space at 6.9~\\kev. In summary, we can find an adequate fit to the data below 4.6~keV with a power-law plus Gaussian emission line model. The fit also seems to prefer absorption greater than that provided by the Galactic column, although this is significant only at the 97\\% confidence level.  Reintroducing the harder data, but leaving the model parameters the same as in Fit \\#6, greatly increases the $\\chi^2$ of the fit (\\#8). The residuals show that there is a clear change in slope at higher energies. Fitting the complete continuum with a broken power-law (\\#9) does a very good job in fitting the entire dataset, and implies that the observed spectrum may be a mixture of two separate emission features. Replacing the broken power-law model with a power-law plus Compton reflection spectrum (with incident power-law included) gives a very similar fit (\\#10; Fig.~\\ref{fig:phy}). \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{Ed151_f2.ps}} \\caption{Best fit model (\\#10) in physical units. The solid line show the total spectrum, while the dashed and dotted lines show the blast wave emission and reflected component, respectively.} \\label{fig:phy} \\end{figure}  The photon-indices of the two power-laws are very different, with the Thomson-thick reflector subject to a very hard ($\\Gamma \\approx 0$) ionizing continuum. The ionization parameter of the slab is well constrained, and, similar to the Gaussian line model, shows that emission from He-like iron is required with strong significance. However, the iron abundance and relative strength of the reflection features cannot be constrained with these data. The above result was obtained assuming a solar abundance of iron and a reflection fraction of unity.    In this paper we have, for the first time, fitted a physically sound model to the spectrum of the X-ray afterglow of a GRB that shows evidence of iron line emission. Our analysis applies to reflection from an optically thick, homogeneous medium, which seems to be required in order to explain the large equivalent width that is observed. The emission feature observed in the X-ray spectrum of GRB~991216 can be explained by reflection if identified with the recombination K$\\alpha$ line from He-like iron at 6.7~keV. Although there exists a significant zone of H-like Fe in the reflector, its K$\\alpha$ line is subject to resonant trapping and is ultimately destroyed by Thomson scattering. If, however, the emitting medium has a velocity dispersion of $\\sim 0.1\\,c$ (Piro et al. 2000), resonant trapping would be suppressed (Lazzati et al. 2001), making the H-like line at 6.97~keV visible. On the other hand, the spectrum can be successfully modelled without kinematic line broadening since the observed width can be reproduced by Comptonization for an expansion velocity below the limit of 0.1$c$ inferred by Piro et al. (2000).  In order to fit the hardening of the spectrum above 5 keV (which was interpreted by Piro et al. (2000) as a recombination edge) we require the ionizing continuum to have a much harder power-law than the observed one. This hard continuum may result from the prompt or central engine emission and would be primarily responsible for ionizing the slab. The other, softer spectrum would then originate from the blast-wave emission, which at later times dominates most of the observed spectrum.  If this power-law spectrum continues to $h\\nu\\ge511$~keV, a significant fraction of the energy in this early ionizing spectrum will be above the $\\gamma\\gamma \\to e^{\\pm}$ formation energy threshold. This will cause new pairs to be formed in the originally optically thick scattering medium, an effect which amplifies the density of scattering charges and increases the temperature of the illuminating material. The effect of pair production on the line luminosity is twofold. On one hand, the electron density is increased, making the recombination time shorter and then the line more luminous. On the other hand, the Thomson opacity of the slab is increased, reducing the depth out of which line photons can escape without being scattered by free electrons. These two effects compensate and as a net result the line luminosity is slightly decreased by the increase of the electron temperature (Kallman et al. 2002).  The results presented here are computed only up to an energy of 100~keV, so we do not take into account pair processes.  We find that iron enrichment is not necessary to reproduce the line strength, although if the true ionization parameter is substantially different than the best-fit value $\\xi \\approx 10^3$, then a moderately super-solar abundance will help bring the line to the required power (Ballantyne \\& Ramirez-Ruiz 2001; Ballantyne et al. 2002). In addition, even if at a non-statistically significant level, the model seems to underestimate the line strength in the grating spectrum, and a supersolar metallicity may cure this. Light elements such as S, Ar and Ca are not currently included in this reflection model (Ross \\& Fabian 1993), so it is not possible to determine if the small deviation at $h\\nu\\sim1.8$~keV is consistent with a S recombination edge.  The value of $\\xi \\approx 10^3$ is predicted by the extreme clumping of ED models, but can also be accommodated in GD models if the reprocessing material is moderately clumped and/or the off-axis burst ionizing continuum is substantially dimmer than the observed burst emission. This is not unexpected, since the beaming break in GRB~991216 occurred quite early (1.5 days; see Halpern et al. 2000). Unfortunately, the statistical quality of the data does not allow clear conclusions to be drawn on the line emission mechanism, and both GD and ED scenarios are consistent with the data.  In summary, our main conclusions are as follows. The iron line-like feature in the X-ray afterglow of GRB~991216 has properties consistent with reflection from an optically thick slab with solar metallicity and ionization parameter $\\xi \\approx 10^3$. The data constrain the energy of the line to be $\\approx 6.8$~\\kev, requiring recombination emission from He-like iron. There is some evidence for substantial intrinsic absorption at the source, consistent with the burst occurring within a gas-rich environment (such as a star-forming region). No kinematic broadening from an outflow is required by the data. The ionizing spectrum that impinges on the slab must however be harder than the afterglow spectrum that dominates the emission at $h\\nu<4$~keV. All these findings can be accounted for by many of the emission mechanisms and geometries discussed in the literature. However, we have shown that ionized reflection models can describe the X-ray afterglow of GRB~991216, and that they may provide important information on the immediate environments of other bursts. Higher signal-to-noise data are required in order to have more insight in the properties and geometry of the line emitting material."
    },
    "0206/astro-ph0206320_arXiv.txt": {
        "abstract": "We announce the initial release of data from the Ohio State University Bright Spiral Galaxy Survey, a $BVRJHK$ imaging survey of a well-defined sample of 205 bright, nearby spiral galaxies.  We present $H$-band morphological classification on the Hubble sequence for the OSU Survey sample.  We compare the $H$-band classification to $B$-band classification from our own images and from standard galaxy catalogs.  Our $B$-band classifications match well with those of the standard catalogs.  On average, galaxies with optical classifications from Sa through Scd appear about one T-type earlier in the $H$-band than in the $B$-band, but with large scatter.  This result does not support recent claims made in the literature that the optical and near-IR morphologies of spiral galaxies are uncorrelated.  We present detailed descriptions of the $H$-band morphologies of our entire sample, as well as $B$- and $H$-band images for a set of 17 galaxies chosen as type examples, and $BRH$ color-composite images of six galaxies chosen to demonstrate the range in morphological variation as a function of wavelength.  Data from the survey are accessible at {\\tt http://www.astronomy.ohio-state.edu/$\\sim$survey/} ",
        "introduction": "Galaxy morphology has been a cornerstone of extragalactic research even before the discovery of the extragalactic nature of the nebulae (e.g., \\markcite{old}Parsons 1850).  Detailed morphological classification of galaxies, in the modern sense dates to the work of \\markcite{tnf}Hubble (1936), who established the now classic tuning-fork diagram for galaxy morphology. Since then there have been a number of modifications of this basic classification scheme.  \\markcite{hag}Sandage (1961) described the sequence in substantially more detail than did \\markcite{tnf}Hubble (1936). \\markcite{dvd}de Vaucouleurs (1959) presents a scheme that accounts for inner ring and spiral structure, as well as strong (SB) and weak (SAB) bar classes. A later version (\\markcite{rc2}de Vaucouleurs, de Vaucouleurs \\& Corwin 1976) introduces the numerical ``T-type'' coding for morphology.  This scheme is described in its most modern form in \\markcite{rc3}de Vaucouleurs et al.~(1991, hereafter RC3).  \\markcite{v76}van den Bergh (1976 and references therein) introduced the DDO luminosity classification scheme that relates the surface brightness and the appearance of spiral arms with the absolute luminosity of the galaxy.  \\markcite{ee2}Elmegreen \\& Elmegreen (1982) created a classification scheme for the degree of structure seen in the spiral arms of galaxies, ranging from flocculent (arm class 0) to grand design (arm class 12). \\markcite{rng}Buta (1986) and \\markcite{rn2}Buta (1995) added a much more detailed classification scheme for resonance rings following the general precepts of \\markcite{dvd}de Vaucouleurs (1959).  The most significant departure from the basic Hubble scheme is that proposed by \\markcite{wwm}Morgan (1958), and given in its final form by \\markcite{mkw}Morgan, Kayser \\& White (1975).  This scheme is essentially a one-dimensional classification based on the degree of concentration of light, with a secondary `form' parameter to distinguish amongst gross galaxy types.  A common property, and, therefore, a common weakness of all these schemes has been the nearly exclusive use of $B$-band plate material for galaxy classification (although see \\markcite{fzw}Zwicky 1955 and \\markcite{sch}Schweizer 1976 for early work on multiwavelength morphology). This constrains existing classification schemes in two major ways.  First, as classification has been done primarily in the $B$-band, it is very sensitive to the distribution of blue stars and dust.  This is particularly unfortunate for the study of late-type galaxies, as the distribution of the young, blue stars and the dust can be very different from the distribution of the total baryonic mass.  Second, although well-exposed plates, taken under good seeing conditions, can offer excellent spatial resolution, plates are poor photometric detectors, and have a very limited dynamic range.  Thus, for instance, it is possible to miss features such as nuclear bars in classification plates that are deeply exposed in order to reveal the structure of spiral arms in the outer disk (see \\markcite{bar}Eskridge et al.~2000 for a discussion of this effect).  The advent of two-dimensional near-infrared (near-IR) detectors has resulted in numerous efforts (e.g., \\markcite{hs}Hackwell \\& Schweizer 1983; \\markcite{hfd}Thronson et al.~1989; \\markcite{bw}Block \\& Wainscoat 1991; \\markcite{bea}Block et al.~1994; \\markcite{bnp}Block \\& Puerari 1999) to compare galaxy morphology in the optical and near-IR.  These studies revealed that there can be substantial differences between the optical and near-IR morphologies of spiral galaxies.  Indeed, \\markcite{bnp}Block \\& Puerari (1999) assert that there is no correlation between the optical and near-IR classifications of spirals.  The main problem with these studies is that they deal with single galaxies, or at most, small samples of objects.  In fact, the lack of a large, well-defined sample of multiwavelength digital imaging of spirals prompted the creation of the OSU Survey.  We believe that classification in both the optical and near-IR, of a large, statistically complete sample of spirals is an important step in our understanding of the physical morphology of galaxies as a function of wavelength.  The astrophysical motivation for such a comparison is that very different sorts of stars dominate the emission from spiral galaxies at optical ($B$-band) and near-IR wavelengths.  In the optical, the regions of current massive star-formation, typically associated with the spiral arms will dominate.  Further, the absorbing effect of interstellar dust is substantial in the optical, and this material is also typically concentrated along the spiral arms.  In the near-IR, the flux is dominated by light from old giants.  There can be significant contributions from young supergiants, but this effect appears to be small in all but the most extreme situations (\\markcite{r98}Rhoads 1998).  A number of recent studies have attempted to classify galaxies in some sort of automated way (e.g., \\markcite{a94}Abraham et al.~1994; \\markcite{oea}Odewahn et al.~1996; \\markcite{nrg}Naim, Ratnatunga \\& Griffiths 1997; \\markcite{onp}Odewahn et al.~2002).  Such efforts, while promising, have mainly been devoted to classification at the very simplest level; determining if a given extended object is a pure spheroid, a disky system, or peculiar.  There have been pioneering efforts aimed at objective bar and spiral structure classification (e.g., \\markcite{a99}Abraham et al.~1999; \\markcite{onp}Odewahn et al.~2002), and such efforts have largely been directed at studying high-redshift galaxy samples.  This is appropriate as such samples are very large (thousands of objects), and each galaxy is typically only a few resolution elements in diameter.  Detailed classification of well-resolved galaxies is still the provenance of ``expert classifiers''.  A landmark study of the classification biases of different ``expert classifiers'' (\\markcite{nea}Naim et al.~1995) reveals that different classifiers are in general agreement with one another, but typically have a scatter of about one T-type.  Whether this is good or bad agreement is, of course, a subjective decision.  We believe it is good enough, given the underlying purpose of morphological classification.  Although descriptive classification is not physical, it is an essential starting point for a physical classification scheme:  Morphology is akin to species taxonomy.  Species taxonomy is not genetics, but it is an essential guide to asking relevant questions in genetics (e.g., \\markcite{bio}Mayr 1942).  In this paper we present the Ohio State University Bright Spiral Galaxy Survey sample; a large, statistically well-defined sample of bright nearby spiral galaxies.  While previous multi-object samples (most notably that of \\markcite{fgt}Frei et al.~1996) have been a great boon for many sorts of projects, none of these samples cover a wavelength range from the optical through the near-IR, and are both large and selected according to a well-defined set of criteria.  We hope that the availability of the OSU Survey data will be of use to the community.  In \\S 2 we describe the survey selection and data taking procedure.  We present statistical results of the comparison of the $B$- and $H$-band morphologies of the sample in \\S 3, and give notes on the $H$-band morphologies in \\S 4.  In \\S 5 we discuss some implications of our results, and suggest promising areas for future research. ",
        "conclusions": "We have presented optical ($B$-band) and NIR ($H$-band) morphological classifications for a statistically complete sample of $\\approx$200 bright spiral galaxies.  Our optical classifications agree well with those of the standard catalogs (see Fig.~2a,b,c and Table 2a,b,c).  The mean absolute difference between our optical classifications and those from both the \\markcite{rc3}RC3 and the \\markcite{cag}CAG is about half a T-type, with no evidence for any systematic differences beyond those already known to exist in the two standard optical catalogs.  Our $H$-band classifications average about 1 T-type earlier than our optical classifications and the optical classifications from both the \\markcite{rc3}RC3 and the \\markcite{cag}CAG (see Fig.~2d,e,f and Table 2d,e,f).  There is thus a clear tendency for spiral galaxies to appear somewhat earlier-type in the NIR than in the optical, but with large scatter.  Our results thus do not support the assertion of \\markcite{bnp}Block \\& Puerari (1999) that the optical and NIR morphologies of spiral galaxies are ``decoupled''.  On average there is quite a good correlation between the optical and NIR morphology of spirals, although there are cases that show dramatic differences.  We have given short descriptions of the $H$-band morphology of our sample.  We have also presented $B$- and $H$-band images of a set of galaxies, selected as examples of their $H$-band morphological type.  In addition, we have presented false-color ($BRH$) images of six galaxies ranging from early- to late-type spirals, three of which look essentially the same in the optical and the NIR, and three of which look substantially different.  The OSU Survey is an important tool for studies of the properties of spiral galaxies in the nearby Universe.  The broad wavelength coverage, and the large, well-defined and statistically complete sample invites a wide range of studies on the morphology, dynamics, ISM, and stellar populations of spiral galaxies. The sample also provides a template to compare with samples of galaxies at high redshift.  Data from the OSU Survey have already contributed to a number of results in studies of galaxy morphology (\\markcite{tea}Terndrup et al.~1994; \\markcite{qea}Quillen et al.~1997; \\markcite{bar}Eskridge et al.~2000; \\markcite{ane}Elmegreen et al.~2002; \\markcite{wyt}Whyte et al.~2002), galaxy dynamics (\\markcite{qfg}Quillen, Frogel \\& Gonzalez 1994; \\markcite{qkd}Quillen et al.~1995; \\markcite{paq}Patsis, Athanassoula \\& Quillen 1997; \\markcite{qnf}Quillen \\& Frogel 1997; \\markcite{pea}Puerari et al.~2000), galaxy evolution (\\markcite{vea}van den Bergh et al.~2002), and dust and extinction (\\markcite{knt}Kuchinski \\& Terndrup 1996; \\markcite{bqp}Berlind et al.~1997; \\markcite{kea}Kuchinski et al.~1998).  We are engaged in a number of collaborative projects using the OSU Survey data. It is particularly important to extend the qualitative study we have done in this paper to more quantitative studies of galaxy morphology.  This work will place our understanding of the properties of current spiral galaxies on a firmer physical footing, and allow for a realistic comparison between the galaxy populations of the low- and high-redshift Universe.  We are engaged in a collaboration to apply the ``Hubble-space'' analysis of \\markcite{anm}Abraham \\& Merrifield (2000) to the OSU sample (\\markcite{vea}van den Bergh et al.~2002; \\markcite{wyt}Whyte et al.~2002).  As the OSU sample is complete and well-defined, it provides for a much better assessment of local galaxy properties than does the sample used by \\markcite{anm}Abraham \\& Merrifield (2000).  The application of the ``Hubble-space'' analysis to our sample allows a better analysis of the claim (\\markcite{a99}Abraham et al.~1999) that the bar fraction of disk galaxies is a strong function of redshift (\\markcite{vea}van den Bergh et al.~2002).  A related project that we are pursuing is a comparative study of a number of the various quantitative measures of ``bar-strength'' that exist in the literature (e.g., \\markcite{ene}Elmegreen \\& Elmegreen 1985; \\markcite{a99}Abraham et al.~1999; \\markcite{bnb}Buta \\& Block 2001; \\markcite{onp}Odewahn et al.~2002).  The pixel-mapping technique of \\markcite{a94}Abraham et al.~(1994) is a powerful tool for studying the star-formation history of nearby galaxies. Application of this technique to a large sample of the OSU survey will result in a more detailed understanding of the relationship between the morphology and star-formation histories of galaxies along the Hubble sequence.  Galaxy asymmetry studies are also proving to be important means of probing the connection between star-formation history and dynamical evolution in disk galaxies (\\markcite{cbj}Conselice, Bershady \\& Jangren 2000).  The OSU survey provides an excellent zero-redshift benchmark for such studies.  It would be especially interesting to apply the methodology of the ``dust-penetrated'' classification scheme of \\markcite{bnp}Block \\& Puerari (1999) to {\\it both} our optical and near-IR data.  We could then carry out a quantitative assessment of the qualitative results of the current study (that spiral galaxies are, on average slightly earlier-type in the NIR that in the optical).  Another related program we are engaged in is a comparison of the optical and IR properties of anemic and normal spirals (\\markcite{ees}Elmegreen \\& Elmegreen 1987; \\markcite{ane}Elmegreen et al.~2002).  Galaxy evolution has become a practical observational study in the last decade.  But the sort of detailed morphological work that is possible for samples of hundreds of galaxies, each hundreds of resolution elements across is simply impossible for samples of tens of thousands of galaxies, each tens of resolution elements across.  The OSU sample provides a zero-redshift benchmark for comparison with high-redshift galaxy samples.  We are involved in collaborations to apply both the ``Hubble-space'' and pixel-mapping techniques of Abraham and collaborators (\\markcite{a94}Abraham et al.~1994; \\markcite{anm}Abraham \\& Merrifield 2000), and the neural-net techniques of Odewahn and collaborators (\\markcite{oea}Odewahn et al.~1996; \\markcite{onp}2002) to the OSU Sample.  It is an irony of modern studies of galaxy evolution that high-redshift galaxy samples are better defined and better observed than zero-redshift samples.  The availability of the OSU Sample helps to rectify this paradoxical situation.  The broad wavelength coverage of the OSU Survey invites extensions to other wavelength regimes.  A number of the galaxies in our sample are part of the BIMA/SONG sample (\\markcite{sng}Regan et al.~2001).  It is a natural question to ask how the distribution of molecular gas in these galaxies correlates with the distribution of both the young stellar populations (as traced by the $B$-band), and the old stellar populations (as traced by the $HK$ bands). Extension of the Survey observations into the vacuum ultra-violet allows both the study of the youngest stellar populations, and provides a firmer footing for comparisons with the rest-frame UV observations of distant galaxies from the Hubble Space Telescope.  We have begun such a program for a few of the galaxies in the OSU Survey (\\markcite{ezo}Eskridge et al.~2001), however, the small field of the of the WFPC2 on HST has restricted this work to the smallest systems in the survey.  The coming availability of the wider-field Advanced Camera for Surveys will allow us to extend the observations to a much larger fraction of the Survey galaxies.  The approaching launch of SIRTF will provide an opportunity to extend the wavelength baseline for the OSU survey through the infrared.  It will be possible to obtain well-resolved images for the OSU sample with IRAC in the range of 3 to 8 microns, and marginally resolved images, and total fluxes for the longer wavelengths with MIPS.  This will allow for a much better understanding of the detailed physical processes relating the stellar populations and the ISM in spirals.  The study of galaxy morphology has been a crucial tool for our understanding of the Universe for the last 80 years.  We are now entering an era when the early qualitative, single wave-band approach to morphology is maturing into a quantitative, multi-wavelength discipline.  Our understanding of the structure, evolution, and underlying physics of galaxies will advance with our ability to study their properties across the electromagnetic spectrum, and along the evolutionary axis of redshift.  We believe the OSU Survey will play an important role in this advance."
    },
    "0206/astro-ph0206050_arXiv.txt": {
        "abstract": "{ In this work we have collected observational data for \\Npsr\\ PSRs. Distances and others parameters for these PSRs were estimated. We present improved distance estimates for radio pulsars by considering importance of their physical properties and improvement of distribution of SFRs (star formation regions) in the Galaxy. For this purpose, both a list of accurate calibrators was constructed and several accurate criteria were established. The following values were calculated from PSRs observational data: luminosities at 400 Mhz and 1400 Mhz, characteristic times, strength of magnetic field and rate of rotation energy. This compilation of data is mainly necessary for statistical investigations and for the physical properties of neutron stars. The whole data is prepared in a publicly accessible web page: \\url{http://www.xrbc.org/pulsar/}. ",
        "introduction": "\\label{s:intro}  It is a well known fact that no relation has been found in pulsar parameters to estimate their distance. For ordinary distant stars, however, one can use the relations either between luminosity and spectral class, or luminosity and pulsation period to estimate their distance. In the early days of pulsar astronomy, since origin of pulsars, mass of their progenitor and their birth rates were not well known, homogeneous electron density distribution was assumed. However, later on pulsars distances have been estimated according to the rough model of Galactic electron distribution and some natural requirements \\citep{% 1981AJ.....86.1953M,1981AZh....58..996G, 1992MNRAS.255..401J,1993ApJ...411..674T, 1996TJPhy..20..275G}. In doing this, one should also know some of the pulsar distances independent to their dispersion measure (DM). 21 cm line of neutral Hydrogen was mainly used in choosing distance calibrators. However, nowadays, calibrators are chosen from members of globular clusters (GCs) or Magellanic clouds (MC), pulsars connected to Supernova remnants (SNRs) with well known distances and from pulsars (PSRs), where their distances are known from other available data.  Irregularities were observed in the distribution of dust, molecular clouds and neutral Hydrogen (HI) in the Galaxy. It is also normal to expect irregularities in electron distribution where the degree of irregularity is (naturally) considerably small. Considerable variations in opacity and polarization can be observed for stars with the same distance in a very small region of sky ($\\sim$1\\degr\\ square) close to the Galactic plane. This is due to a very inhomogeneous distribution of dust clouds. For Hydrogen column density along the line of sight there are two surveys where they studied large number of stars; one with 554 stars \\citep{1994ApJS...93..211D} and the other with 594 stars \\citep{1994ApJS...94..127F}. They both show that irregularities in HI distribution are quite different than the dust and molecular cloud distribution \\citep{1998A&AT...17..301A}. The dispersion measure (DM), which is connected with the electron distribution, changes also for pulsars of similar distances, and for close regions of the sky. These irregularities in electron distribution are due to contribution of both HII regions and SNRs along the line of sight, and gravitational potential and gas temperature distribution in the Galaxy. But irregularities in electron distribution is considerably smaller than the ones in other components of interstellar medium that we have mentioned above. Even though these irregularities are small, there is no simple model for Galactic electron distribution to calculate each pulsar's distance. Moreover, constructing a complex model which requires a lot of data for interstellar medium and PSRs, \\citep[\\eg][]{1993ApJ...411..674T} cannot avoid large errors for individual pulsars.  In order to investigate the arm structure around the Sun within a distance of 4--5 kpc, usually objects like OB associations and open clusters (OC) are studied. For these objects the relative errors in estimating their distances could reach 30\\% \\citep{% 1978ApJS...38..309H,1989QB806.E37......, 1992A&AS...94..211G,1995A&AS..109..375A}. There is no single good method to estimate the distance of all extended objects belonging to the arms (molecular clouds, neutral Hydrogen clouds, SNRs and HII regions). In determining the distances to these objects using HI 21 cm line and Galaxy rotation models, error exceeds 30\\% and it increases with distance and in the vicinity of longitudes $l=0$\\degr\\ and $l=180$\\degr. However, it is the most widely used model. For distant X-ray sources, Hydrogen column density is used as another method in estimating distances. However, error in this method is also large. Since progenitors of pulsars are massive stars, their birth places are in the star formation regions (SFRs). Furthermore, even though young pulsars with characteristic age of $\\tau<$\\EE{5}{5} years have high space velocities, they cannot escape from their birth places. Thus, if number of young pulsars discovered increases and distances to these pulsars are well known then farther away arm structures could be studied.  Archiving radio pulsar data, dates back to 1981. The first full catalog included 333 pulsars which covered discoveries up to 1980 \\citep{1981AJ.....86.1953M}. The next catalog which plays an important role in pulsar astronomy contained 706 pulsars \\citep{1996unpubwork.....T}. This one covered both old (since 1981) and new pulsars \\citep[][and some others]{% 1985ApJ...294L..25D,1985Natur.317..787S, 1986ApJ...311..694S,1992MNRAS.254..177C, 1992MNRAS.255..401J,1993ApJ...411..674T}. This last catalog has not been updated since then. However, individual pulsars can be reached through a publicly accessible web page \\footnote{\\small\\url{http://pulsar.ucolick.org/cog/pulsars/catalog/}}. Since 1996, several pulsar surveys have been carried out \\citep{% 1995MNRAS.274L..43J,1996MNRAS.279.1235M, 1997ApJ...478L..95S,1998MNRAS.295..743L, 2000MNRAS.312..698L,2001ApJ...548L.187C, 2001ApJ...547L..37E,2001ApJ...553..801E, 2001MNRAS.326..358E,2001ApJ...560..365E, 2001PASA...18....1M}. In addition to this, inner regions of SNRs have been scanned to search for pulsars with connections to SNRs \\citep{% 1996ApJ...458..257G,1998A&A...331.1002L,1996AJ....111.2028K}. After 1996, the following pulsars with connections to SNRs or pulsars with confirmed association connections have been found and their distances were accurately determined (see Table \\ref{t:calib}):\\\\ \\begin{tabular}{@{}l l@{}} J0205+64/G130.7+3.1   &\\citep{2002ApJ...568..226M},\\\\ J1119-6127/G292.2-0.5 &\\citep{2001ApJ...554..152C},\\\\ &\\citep{2001ApJ...554..161P},\\\\ J1124-5916/G292.0+1.8 &\\citep{2002ApJ...567L..71C},\\\\ J1803-2137/G8.7-0.1?  &\\citep{1994ApJ...434L..25F},\\\\ J1846-0258/G29.7-0.3  &\\citep{2000ApJ...542L..37G},\\\\ J1952+3525/G69.0+2.7  &\\citep{1990ApJ...364..178K},\\\\ J2229+6114/G106.6+2.9 &\\citep{2001ApJ...552L.125H}.\\\\ \\end{tabular}\\\\ Furthermore, Globular Clusters (GC) have been also searched for pulsars \\citep{% 1995mpds.conf...35L,1996IAUS..174..181K, 1996MNRAS.282..691B, 2000ApJ...535..975C,2001ApJ...548L.171D}. In globular cluster NGC104 (47 Tuc) 10 pulsar up to 1996 and 10 more pulsar after 1996 have been found. For the other known globular clusters no new pulsars were found. However, in each globular clusters NGC 6266, NGC 6342, NGC 6397, NGC 6544 and NGC 6752 one pulsar has been found after 1996 (Table \\ref{t:calib}). \\input Tcalib.tex  In early days of pulsar observations a base frequency of around 400 MHz was used in the search. Since DM values of distant pulsars are high, 1400 MHz was used in surveys and in search for PSRs in SNRs and GC. As expected, the newly discovered pulsars are generally in the direction of the Galactic center. After 1996, no new pulsars have been found in Magellan Clouds (MC). However, number of pulsars in GCs and number of millisecond pulsars with known ages ($P <$ 0.1 sec and \\pdot $<$ \\EE{}{-16} sec/sec) increased about 1.5 and 1.4 times, respectively. There is an considerable increase in number of pulsars found with low fluxes due to increase in both sensitivity of instrumentation used in pulsar surveys and the number of detailed surveys. For example in Arecibo's survey window (40\\degr$\\le l\\le$65\\degr; $|b|\\le 2.5$\\degr) 12 new pulsars were found. In this article our aim is to combine both old and new observational pulsar data and to calculate their parameters. ",
        "conclusions": ""
    },
    "0206/astro-ph0206099_arXiv.txt": {
        "abstract": "{ We study the statistical characteristics of a box-fitting algorithm to analyze stellar photometric time series in the search for periodic transits by extrasolar planets. The algorithm searches for signals characterized by a periodic alternation between two discrete levels, with much less time spent at the lower level. We present numerical as well as analytical results to predict the possible detection significance at various signal parameters. It is shown that the crucial parameter is the effective signal-to-noise ratio --- the expected depth of the transit divided by the standard deviation of the measured photometric average within the transit. When this parameter exceeds the value of 6 we can expect a significant detection of the transit. We show that the box-fitting algorithm performs better than other methods available in the astronomical literature, especially for low signal-to-noise ratios. ",
        "introduction": "A considerable fraction of the periodic astronomical time series can be modeled rather accurately by finite sums of sinusoidal components. In general, these Fourier-sums have a single dominant component, and therefore the basic method of Discrete Fourier Transformation (DFT) has become commonplace in almost all applications (e.g., Deeming 1975). When the signal becomes distorted by higher harmonics (e.g., light curves of fundamental mode RR~Lyrae and $\\delta$~Cephei stars), this simple approach fails to perform properly, due to leakage of the signal power to many higher harmonics. One way to deal with this problem is to use a multifrequency Fourier fit for better approximation of the signal shape and thereby to increase the algorithm efficiency, a method recently suggested by Defa\\\"y, Deleuil \\& Barge (2001) for the search for extrasolar planetary transits. Another generally accepted approach is the so-called Phase Dispersion Minimization (PDM), which searches for the best period that yields the `smoothest' folded time series.  Application of variants of the PDM method in the analyses of variable star observations goes back to earlier times than that of the DFT. This is primarily because the PDM algorithm does not require the computation of trigonometric functions, which put a heavy load on the computers, especially in those early days. The most frequently cited implementation of the PDM idea is that of Stellingwerf (1978). However, earlier versions had appeared already in the '60s and early '70s, like those of Lafler \\& Kinman (1965, hereafter the L-K method), Jurkevich (1971) and Warner \\& Robinson (1972, hereafter the W-R method). Actually, it can be shown that up to a frequency- (or trial period-) independent constant, the method of Jurkevich (1971) is equivalent to that of W-R (see, Kov\\'acs 1980). Furthermore, without the additional feature of overlapping bin structure, the method of Stellingwerf (1978) is equivalent to that of Jurkevich (1971).  The study of the algorithm presented in this paper has been stimulated by the increasing interest in searching for periodic transits by extrasolar planets (e.g., Gilliland et al.  2000; Brown \\& Charbonneau 2000; Udalski et al.\\ 2002), which follow the discovery of the transit of HD~209458 (Charbonneau et al.\\ 2000; Henry et al.\\ 2000) by its planetary companion (Mazeh et al.\\ 2000).  Due to the short duration of the transit relative to the orbital period (typically less than 5\\%), the signal expected is extremely non-sinusoidal. Considering the shallowness of the transit (typically less than 2\\% in the case of Jupiter-size planets) and the expected high noise level of the ground-based small telescopes capable of performing large-scale surveys (e.g., Borucki et al. 2001), we suggest an algorithm that utilizes the special form of the signal. The algorithm studied here (see also Gilliland et al. 2000 and Udalski et al.\\ 2002) is based on direct Least Squares (LS) fits of step functions to the folded signal corresponding to various trial periods. We present numerical simulations as well as analytical considerations to estimate the ability of the algorithm to detect a faint signal in a noisy time series. We show that the algorithm performs significantly better and more efficiently than the published variants of PDM, DFT or some LS modification of the latter. ",
        "conclusions": "This paper has examined the statistical characteristics of the Box-fitting Least Squares algorithm to detect periodic transits in time series of stellar photometric observations. The algorithm strongly relies on the anticipated box-shape of the periodic light curve. The advantage of using a predetermined shape of the light curve manifests itself in the high efficiency of this method relative to the other search methods, which are generic and can detect any periodic variation.  The algorithm studied here assumes only two levels of the periodic light curve. This assumption ignores all other features that are expected to appear in planetary transits. Thus, we ignore the gradual ingress and egress phases of the transit, which carry important information about the parameters of the planetary orbit (e.g, Sackett 1999). The lengths of these phases are short compared to the transit and thus they are not expected to affect significantly the results of the search.  Another effect we ignore is the limb-darkening effect, which has indeed been shown to be small in the case of HD~209458 (e.g., Deeg et al.\\ 2001).  The effectiveness of the algorithm relies on the above simplifying assumption, which is justified as long as we are interested in a detection tool. After the periodicity is detected we can try to recover subtle features of the folded light curve, in order to derive the stellar and the planetary characteristics.  Our main interest is in cases where the signal-to-noise ratio is small, and one cannot identify the signal by monitoring a single transit, because the stellar drop in intensity is buried in the noise. Contrary to the search of transits by the HST in 47 Tuc (Gilliland et al.\\ 2000), where the noise was small relative to the expected transit dip, we concentrate on cases in which the periodic signal can be detected only after many measurements are accumulated and the unknown transit is monitored many times. To be able to deal with a large number of observations, of a thousand or more, we have introduced binning into the folded data. We have shown that as long as the bin size is small compared to the expected transit length, the efficiency of the method is not affected.  One additional factor that determines the computational load of the algorithm is the range of transit length searched for. The maximum possible transit length can be estimated if we know the orbital, stellar and planetary radii. For a given stellar mass, the stellar radius can be derived by the mass-radius relation, and the orbital radius can be derived for any period. Recent theories give some estimates for the planetary radii. Therefore, for a given stellar mass we can estimate the maximum duration of the transit, which for HD~209456 is only a few percent of the period. For most ground-based and space searches for planetary transits one would have some idea of the stellar mass of all transit candidates, and therefore we can make our algorithm computationally more efficient by imposing a variable maximum duration on the transit length.  The significance of the detection depends primarily on the effective signal-to-noise ratio of the transit. The signal is the stellar brightness within the transit, relative to the brightness outside the transit, and the noise is the expected scatter of the measured average of the stellar brightness inside the transit. The scatter is composed, obviously, of the observational noise as well as of stochastic variation of the stellar intensity. It seems that the effective signal-to-noise ratio should exceed $6$ in order to get a significant detection. This requirement should be taken into account when planning future searches for extrasolar planetary transits."
    },
    "0206/hep-ph0206092_arXiv.txt": {
        "abstract": "Neutralinos are natural candidates for cold dark matter in many realizations of supersymmetry. We briefly review our recent results in the evaluation of neutralino relic abundance and direct detection rates in a class of supergravity models. \\vspace{1pc} ",
        "introduction": "The fact that most of the Universe is dark and composed of new, exotic components has been established over the years by means of various sets of observations, which recently have converged into a consistent picture where the Universe is about critical and dominated by an exotic form of matter and by an unexpected type of dark energy. In terms of the density parameter $\\Omega$, the current view can be summarized as follows: the total amount of matter/energy of the Universe is $\\Omega_{\\rm tot} \\simeq 1$ at the 10\\% level and this is composed of a matter component $\\Omega_{\\rm M} \\simeq 0.3$ and a vacuum--energy component $\\Omega_{\\Lambda} \\simeq 0.7$ \\cite{TAUP01_Proc}. The existence of both dark exotic matter and dark energy asks for extension of the standard model of fundamental interactions, since no known particle or field can represent either of these components. In this paper, we will deal with the problem of explaining the observed amount of dark matter, which we can summarize as: $0.05 \\lsim \\Omega_{\\rm M} h^2 \\lsim 0.3$, and with the studies related to the searches for dark matter particles. For an updated review on these subjects, see Ref. \\cite{TAUP01}. ",
        "conclusions": "We can certainly define the following items as the current main issues and open problems in particle dark matter studies: {\\em i)} to explain the observed amount of dark matter in the Universe ($0.05 \\lsim \\Omega_{\\rm M} h^2 \\lsim 0.3$) by finding suitable particle candidates; {\\em ii)} to detect a relic particle. We have shown that for both of these issues, there appear to be good prospects of success, especially for the most studied candidate which is the neutralino. In particular, there are many susy schemes where relic neutralinos can provide enough cosmological abundance to explain the observed amount of dark matter, and at the same time they can have detection rates large enough to be accessible to detection. Clearly the occurrence of this particularly interesting situation depends on the actual realization of supersymmetry. The observation of a signal from dark matter, like for instance in the case of the annual modulation effect observed by the DAMA/NaI Collaboration or of signals which could hopefully come in future experiments, can be very important not only for astrophysics and cosmology but also for particle physics, since the need to explain such effects can help in deriving properties of particle physics models and possibly discriminate among different realizations, for instance of supersymmetry."
    },
    "0206/astro-ph0206272_arXiv.txt": {
        "abstract": "#1{\\vskip 0.5cm {#1}} \\def\\msol{{\\cal M}_\\odot} \\def\\lsol{{\\cal L}_{\\odot,B}} \\def\\lso{{\\cal L}_{\\odot}} \\def \\refs {\\begingroup \\frenchspacing \\parindent = 0 pt \\everypar = {\\hangindent = 20.0 pt \\hangafter = 1}} \\def \\endrefs {\\par \\endgroup} \\input psfig.tex \\begin{document}   \\title{Cluster-Group Interaction in the Virgo Cluster}  \\authors{S. Schindler$^{}$}  \\affiliation{ University of Innsbruck, Institute for Astrophysics, Technikerstr. 25, 6020 Innsbruck, Austria}  \\abstract{We present two projects related to interaction in the Virgo cluster. In the first section we draw a quantitative comparison of the distribution of the galaxies and the intra-cluster gas taking into account that the Virgo cluster has an irregular structure consisting of several subclusters. In the second section we show hydrodynamic simulations of the interaction (ram pressure stripping) of a galaxy like M86 with the intra-cluster gas.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206044_arXiv.txt": {
        "abstract": "We present numerical simulations of the evolution of low-mass, isothermal, molecular cores which are subjected to an increase in external pressure $P\\xt$. If $P\\xt$ increases very slowly, the core approaches instability quite quasistatically. However, for larger (but still quite modest) $dP\\xt/dt$ a compression wave is driven into the core, thereby triggering collapse from the outside in. If collapse of a core is induced by increasing $P\\xt$, this has a number of interesting consequences. (i) The density profile is approximately flat in the centre during the prestellar phase (i.e. before the compression wave converges on the centre creating the central protostar). (ii) During the prestellar phase there are (subsonic) inward velocities in the outer layers of the core, whilst the inner parts are still approximately at rest. (iii) There is an initial short phase of rapid accretion (notionally the Class 0 phase), followed by a longer phase of slower accretion (the Class I phase). All these features accord well with observation, but are at variance with the predictions of the standard theory of star formation based on the inside-out collapse of a singular isothermal sphere. We note that the setting up of a coherent inward velocity field appears to be a generic feature of compression waves; and we speculate that interactions and interference between such velocity fields may play a crucial r\\^ole in initiating the fragmentation of cores and the genesis of multiple star systems. ",
        "introduction": "The initial conditions for star formation are in general rather weakly constrained. The observations (e.g. Andr\\'e, Ward-Thompson, \\& Barsony 2000; Myers, Evans, \\& Ohashi 2000) are limited by telescope resolution and confusion, and their interpretation is hindered by the complexities and uncertainties of gas-phase abundances, excitation conditions, and radiation transport. However, a picture is emerging which suggests that star formation is sometimes triggered rather impulsively.  Molecular-line mapping by Myers and collaborators (e.g. Myers \\& Benson 1983; Benson \\& Myers 1989) and subsequent searches by Clemens \\& Barvainis (1988), Bourke, Hyland \\& Robinson (1995), Jessop \\& Ward-Thompson (2000), have established a large sample of dense cores which appear to be the sites of ongoing or imminent star formation. Many of these cores contain IRAS sources (Beichman et al. 1986), and are therefore presumed to have already formed protostars. Those which do not contain IRAS sources are termed starless cores.  Submillimeter mapping of these starless cores by Ward-Thompson et al. (1994, 1999) has identified a subset which, on the basis of their virial ratios and relatively high central densities, are likely to be in a state of imminent or on-going contraction. These starless cores are therefore termed prestellar (strictly pre-protostellar). Submillimeter mapping of prestellar cores (Ward-Thompson et al. 1994, Andr\\'e et al. 1996, Ward-Thompson et al. 1999) suggests that, if the emitting dust is isothermal, the cores have rather flat density profiles in the centre. Specifically, $\\eta \\equiv - d \\ell n [\\rho] / d \\ell n [r]$ is in the range 0 to 1, in the innermost few thousand AU. This conclusion may be somewhat weaker if, as seems likely (Jessop \\& Ward-Thompson 2001; Evans et al. 2001; Zucconi, Walmsley \\& Galli 2001; Ward-Thompson, Andr\\'e \\& Kirk, 2001), the dust temperature decreases towards the centre of a core. However, mid-infrared absorption measurements of prestellar cores -- which are insensitive to the temperature profile -- also indicate flat inner density profiles (Bacmann et al. 2000). Additionally, mid-infrared observations suggest that beyond $\\sim$ 10,000 AU the density profile may steepen to $\\eta \\ga 4\\,$ (Abergel et al. 1996, Bacmann et al. 2000). This steepening of the outer envelope appears to occur at smaller radii in relatively close-packed protoclusters like $\\rho$ Ophiuchi than in regions of distributed star formation such as Taurus; viz. at $\\sim$ 3,000 AU in $\\rho$ Ophiuchi, and at $\\sim$ 15,000 AU in Taurus (Motte \\& Andr\\'e 2001).  The asymmetric self-absorbed molecular-line profiles observed in some prestellar cores (Tafalla et al. 1998, Williams et al. 1999, Lee, Myers \\& Tafalla 1999, Gregersen \\& Evans 2000) suggest that they are indeed already collapsing. In particular, the detailed analyses of L1544 by Tafalla et al. (1998) and Williams et al. (1999) imply that the inner parts of the core are relatively stationary, and an approximately uniform velocity field has been established in the outer layers. This is very reminiscent of the velocity fields which are set up by inward-propagating compression waves in similarity solutions for contracting isothermal spheres (Whitworth \\& Summers 1985). It is this apparent similarity which we explore in the present paper.  The prestellar phase terminates as soon as a star-like object forms at the centre of the dense core, and the core then becomes a Class 0 protostar. Conceptually, the Class 0 protostellar phase terminates once the extended envelope contains less than half the total mass of the original core, i.e. more than half the mass is in the central star-like object(s) plus attendant disc(s) (Andr\\'e, Ward-Thompson \\& Barsony 1993, 2000); the Class I phase then begins (Lada \\& Wilking 1984, Lada 1987, Andr\\'e \\& Montmerle 1994). The Class I protostellar phase terminates when most of the envelope has been accreted or dissipated, revealing a classical T Tauri star (CTTS) accreting from a residual circumstellar disc, i.e. a Class II object. Once the inner disc has been dissipated, the accretion rate is greatly reduced and the source becomes a weak-lined T Tauri star (WTTS) or Class III object.  On the basis of statistical arguments (i.e. source numbers and a presumed constant star formation rate) it is inferred (Beichmann et al. 1986, Andr\\'e, Ward-Thompson \\& Barsony 2000) that the prestellar phase lasts, typically, $10^6$ to $10^7$ years. In close-packed protoclusters like $\\rho$ Ophiuchi, the Class 0 phase appears to last a few times $10^4$ years, and is characterized by powerful collimated outflows indicative of rapid accretion, $\\ga 10^{-5} \\, M_\\odot \\, \\mbox{year}^{-1}$ (Bontemps et al. 1996). In distributed star formation regions like Taurus, the duration of the Class~0 phase appears to be longer, $\\sim 10^5$ years, and the accretion rates lower, $\\ga 2 \\times 10^{-6} \\, M_\\odot \\, \\mbox{year}^{-1}$ (Motte \\& Andr\\'e 2001). The Class I phase appears to last $\\sim 2 \\times 10^5$ years (e.g. Greene et al. 1994, Kenyon \\& Hartmann 1995) and is characterized by slower accretion, $\\la 10^{-6} \\, M_\\odot \\, \\mbox{year}^{-1}$, and weaker less collimated outflows (Bontemps et al. 1996). Together, the Class II and Class III phases appear to last $\\ga 10^7$ years, but seemingly the transition from Class II (CTTS) to Class III (WTTS) is very short and can occur at any time; there are both very young WTTSs (close to the birthline, on the right of the Hertzsprung-Russell Diagram) and very old CTTSs (approaching the Main Sequence on the left of the Hertzsprung-Russell Diagram) (e.g. Stahler \\& Walter 1993).  These observationally inferred features of prestellar and protostellar evolution combine to form a reasonably coherent picture. However, especially in star-forming clusters, this picture is difficult to reconcile with the standard theory of Shu, Adams \\& Lizano (1987) based on the inside-out collapse of a singular isothermal sphere. The strength of the standard theory is that it makes specific quantitative predictions, but some of these predictions are difficult to reconcile with observation. (i) In the standard model, prestellar cores should be strongly centrally condensed, $\\eta \\equiv - d \\ell n [\\rho] / d \\ell n [r] \\sim 2$ and young protostars should be somewhat less centrally condensed, $\\eta \\sim 3/2\\,$. By contrast, observations suggest that prestellar cores have rather flat central density profiles, $\\eta \\la 1$, and Class 0 protostars have steeper ones. (ii) The standard theory predicts that prestellar cores are static, and that inward motions only develop during the protostellar phase and are initially confined to the central regions. In contrast, the observations imply that inward motions already exist during the prestellar phase, and that initially they are more rapid in the outer regions. (iii) The standard theory predicts that the accretion rate is roughly constant, and hence the Class 0 and Class I lifetimes should be comparable. Observations detect many more Class I sources than Class 0 ones -- although this statistical result may be compromised by the difficulty of measuring precisely when an object has accreted half the total mass of its initial core. (iv) Additionally, the standard theory assumes initial conditions which are unlikely to arise in nature, because they are both singular and unstable. (v) Also the standard theory has a strong inbuilt pre-disposition to the formation of single stars -- in stark contrast with the high proportion of binaries and higher multiples observed in young star-formation regions. Therefore, although the standard theory may provide a good, zeroth-order description of protostellar collapse in sparse, quiescent  star-formation regions like Taurus (cf. Motte \\& Andr\\'e 2001), it appears that more dynamical models are required to understand close-packed regions like $\\rho$ Ophiuchi.  Foster \\& Chevalier (1993) have explored how the collapse of an isothermal core develops if one abandons the assumption of singularity. Their simulations start from non-singular isothermal equilibria, i.e. Bonnor-Ebert spheres (e.g. Bonnor 1956),  and collapse is then triggered by discontinuously increasing the density. The ensuing collapse results in supersonic inflow velocities in the centre of the core at the end of the prestellar phase, and in a marked decline of the accretion rate from the Class 0 to the Class I phase (see Henriksen, Andr\\'e, Bontemps 1997 and Whitworth \\& Ward-Thompson 2001 for simpler, pressure-free descriptions of this evolution).  In this paper we pursue the consequences of non-singularity further. Our simulations also start from non-singular isothermal equilibrium cores, but collapse is then triggered by a steady increase in the external pressure (cf. Myers \\& Lazarian 1998). Most of the observational constraints detailed above can be reproduced rather well by this model.  In Section 2 we describe the numerical method we use, and the initial and boundary conditions. In Section 3 we present the results, and in Section 4 we discuss them. Section 5 summarizes our main conclusions. ",
        "conclusions": "We conclude that the model of protostellar collapse developed here merits further investigation. It appears to reproduce the gross features of the density and velocity fields observed in prestellar cores and protostars, and the relative ages of the Class 0 and Class I phases. We now need to explore the consequences of introducing an energy equation, the effect of initial rotation, and the influence of perturbations as a means of inducing the formation of binaries and higher multiples."
    },
    "0206/astro-ph0206334_arXiv.txt": {
        "abstract": "We have observed the narrow--line regions (NLRs) of the seven brightest radio--quiet PG (or BQS) quasars (z $<$ 0.5) with the Wide Field and Planetary Camera 2 on board the Hubble Space Telescope (HST). Linear--ramp filters were used to image the [\\ion{O}{3}]\\,$\\lambda$5007 line emission with 0\\farcs0455--0\\farcs1 pixel resolution. We find that the NLRs are very compact with typical extents of 2\\arcsec--4\\arcsec. Two quasars show compact filamentary structures similar to Seyfert NLRs. They may be related to radio outflows.  Most interestingly, when including a sample of Seyfert galaxies observed with HST, we tentatively find that the size of the NLR is proportional to the square root of the [\\ion{O}{3}] luminosity. This is comparable to the scaling found for the size of the broad--line region with continuum luminosity, which has been interpreted in terms of a constant photoionization parameter. The relation determined here connects the NLR of radio--quiet quasars and Seyferts over three orders of magnitude in [\\ion{O}{3}] luminosity. ",
        "introduction": "Quasars are active galactic nuclei (AGN) in which two different regions of ionized gas can be distinguished --- the broad--line region (BLR) and the narrow--line region (NLR) which exhibit optical emission lines, kinematically broadened with typical widths of 10$^{3-4}$\\,km\\,s$^{-1}$ and 10$^{2-3}$\\,km\\,s$^{-1}$, respectively.  The optical spectra of quasars resemble those of the less luminous Seyfert galaxies, and it is presumed that most radio--quiet quasars are their more luminous counterparts. The most prominent optical emission lines of the NLR are [\\ion{O}{3}]\\,$\\lambda$5007 (hereafter referred to simply as [\\ion{O}{3}]) and H$\\alpha$\\,+\\,[\\ion{N}{2}] $\\lambda\\lambda$6548,6584.  In the Unified Model of AGN, an optically thick obscuring dust torus is envisioned to encircle the accretion disk \\citep{ant93}, leading to a so--called ``ionization cone'' formed by anisotropic escape of ionizing photons \\citep{sch88,pog88,pog89,sto92}. Thus, the NLR is expected to show a specific morphology, which is suited for investigation with the high spatial resolution afforded by the Hubble Space Telescope (HST). Such investigations have been carried out for Seyfert galaxies, revealing highly elongated structures or ionization cones \\citep{cap96,fal98}.  While the NLR in Seyferts is now relatively well studied, there are no comparable studies for the NLR in quasars. By observing quasars we can greatly extend the luminosity range to answer questions such as: Are quasars indeed just scaled--up versions of Seyferts?  Does the size of the NLR scale with luminosity? Is this emission--line region possibly affected by radio jets shaping the interstellar medium, or is there a contribution from star formation?  We make a first step in this direction by presenting an HST emission--line imaging survey of a complete sample of the seven brightest [in [\\ion{O}{3}], \\citet{bor92}] radio--quiet quasars from the BQS \\citep{sch83,kel89} with z $<$ 0.5. The sample is given in Table~\\ref{tabobs}. Here we present an investigation of the NLR structure as seen in the [\\ion{O}{3}] line and discuss the size of the NLR as a function of luminosity.  Luminosity distances for both Seyferts and quasars were calculated by using redshifts relative to the 3K background as derived with the velocity calculator provided by the NASA Extragalactic Database (NED). Throughout this paper we adopt a Hubble constant of $H_0$ = 65 km\\,s$^{-1}$\\, Mpc$^{-1}$ and a homogeneous, isotropic, flat world model, which includes Einstein's cosmological constant $\\Lambda$ in agreement with recent supernova measurements \\citep{per99}: $\\Sigma \\Omega$ = 1, $\\Omega_{{\\rm matter}}$ = 0.35, $\\Omega_{{\\rm radiation}}$ = 0.05, and $\\Omega_{\\Lambda}$ = 0.6. ",
        "conclusions": "We find that the NLRs of seven bright radio--quiet PG quasars are remarkably compact with typical extents of 2\\arcsec--4\\arcsec. Hence, detailed imaging of quasar NLRs requires sub--arcsecond resolution. Generally, the structure is relatively symmetric, in agreement with the unified scheme, that predicts a view into the ionization cones of these type 1 objects.  Two quasars exhibit compact filamentary structure like that seen for Seyferts. These structures may be related to radio outflows. This is reminiscent of the situation in many Seyfert galaxies, where radio outflows are morphologically related to the NLR. In three quasars the radio emission is approximately as compact as the NLR, and in all five cases that can be found in literature, the maximum angular extent is in agreement with our compact extents.  The NLR in quasars seems to be consistent with being a scaled--up version of the NLR in Seyferts. In fact, one of the Seyfert 2 galaxies, \\objectname{Mrk 34}, has a similar radius and [OIII] luminosity to the quasars, and hence might be considered a ``type 2 quasar''. In addition, the size of the NLR seems to scale roughly with the square root of the [\\ion{O}{3}] luminosity when combining the quasar and Seyfert samples.  The latter result, however, has to be taken with some caution. For one, we have a very limited number of sources available. Especially for quasars, high resolution images of the NLR are extremely scarce, one essentially needs the resolution of HST plus the flexibility of the LRFs to image redshifted emission lines. Second, when comparing the Seyfert 2 and the quasar sample, one needs to consider that, according to the unified scheme, orientation effects might bias the type 2 sample towards somewhat larger sizes. Third, the individual scatter of NLR sizes is relatively large. Fourth, the ``size'' of the NLR is not a well defined quantity, being dependent on sensitivity and resolution. Nevertheless, given the large span in luminosities, none of the effects is likely to change the results significantly.  Seeking a clue to the origin of this relation, we consider a recombination line like H$\\beta$, rather than [\\ion{O}{3}], as a tracer of continuum luminosity. The ionization parameter is given by $U = Q/(4 \\pi c n_e R_u^2)$ ($Q =$ rate of H--ionizing photons, $R_u$ = distance between photoionizing source and emission--line clouds). Employing the relation $\\omega Q = \\alpha_B/(\\alpha_{{\\rm H}\\beta} h \\nu_{{\\rm H}\\beta}) = 3.9\\,10^{12} L_{{\\rm H}\\beta}$ (cgs units; $\\omega$ is the covering factor and the recombination coefficients $\\alpha_B$ and $\\alpha_{{\\rm H}\\beta}$ are taken from \\citet{ost89} for $T=2\\,10^4$K) yields a relation $R_u \\propto L_{{\\rm H}\\beta}^{0.5}$ for given $U$, $n_e$, and constant $\\omega$.  The slope in Figure~\\ref{o3corr} is close to the slope of $0.6\\pm0.1$~given by \\citet{pet01} for the relationship between BLR size (measured from reverberation mapping) and the continuum--luminosity (at 5100\\AA). \\citet{mcl02} find an even tighter correlation between BLR radius and 3000\\AA~luminosity, with both correlations being consistent with a relation of the form $R_{\\rm {BLR}}\\propto L^{0.5}$.  To first order, BLRs and NLRs exhibit similar ionization parameters commonly explained by interaction of clouds with a quasar wind \\citep{sch86}. More recently, \\citet{dop02} use radiation--pressure dominated photoionization models to explain a constant ionization parameter (in the range of --2.5 $\\le$ $\\log U$ $\\le$ --2).  At the outskirts of a NLR, $n_e\\sim 10^{2-3}$ cm$^{-3}$, $U\\sim 10^{-(2-3)}$, so that efficient [\\ion{O}{3}] emission comes from regions with $U n_e \\approx 1$, corresponding to an ionizing photon rate $Q/(4 \\pi R_u^2) = 3~10^{10}$ photons/(s cm$^2$).  Individual values of $R_u$ derived with this scheme are slightly larger than the measured $R$, corresponding to covering fractions $\\omega$ $\\le$ 1.  The same average [\\ion{O}{3}]/H$\\beta$ ratio for Seyferts and quasars preserves the slope of the size--[\\ion{O}{3}] relation for a size--H$\\beta$ relation. This corroborates the explanation in terms of the ionization parameter and argues for photoionization as the main underlying process for producing narrow lines.  As pointed out by the referee, the relationship $L_{\\rm [OIII]}$ (or $L_{H\\beta}$) $\\propto R_{\\rm NLR}^2$ corresponds to constant [\\ion{O}{3}] (or H$\\beta$) surface brightness. We are, however, not aware of any other physical, or instrumental, reason why this should be the case, since [\\ion{O}{3}] is emitted by the whole {\\em volume} filled with NLR clouds.  Clearly, more HST observations of quasars with lower [\\ion{O}{3}] luminosity are needed to test the validity of the correlation.  We expect fainter [\\ion{O}{3}] quasars to have even smaller NLRs than the sources presented here and such a program will present an observational challenge for the years to come. Nevertheless, the available data provides the first direct evidence that quasar and Seyfert NLR are related and possibly evolve along a common luminosity--size track determined by photoionization."
    },
    "0206/astro-ph0206102_arXiv.txt": {
        "abstract": "We present a detailed comparison between predicted and empirical $PL_{I,K}$ relations and Wesenheit function for Galactic and Magellanic Clouds (MCs) First Overtone (FO) Cepheids. We find that zero-points predicted by Galactic Cepheid models based on a noncanonical (mild overshooting) Mass-Luminosity (ML) relation are in very good agreement with empirical zero-points based on HIPPARCOS parallaxes, while those based on canonical (no overshooting) ML relation are $\\approx 0.2-0.3$ mag brighter. We also find that predicted and empirical $PL_K$ relation and Wesenheit function give, according to optical ($V,I$ OGLE) and near-infrared (NIR, $K$, {\\sc 2mass}) data, mean distances to the MCs that agree at the 2\\% level. Individual distances to the Large and the Small Cloud are: $18.53\\pm0.08$-$19.04\\pm0.11$ (theory) and $18.48\\pm0.13$-$19.01\\pm0.13$ (empirical). Moreover, predicted and empirical FO relations do not present, within the errors, a metallicity dependence. Finaly, we find that the upper limit in the FO period distribution is a robust observable to constrain the accuracy of pulsation models. Current models agree within 0.1 in $\\log P$ with the observed FO upper limits. ",
        "introduction": "The massive photometric databases collected by the micro-lensing experiments (MACHO, EROS, OGLE) substantially increased the number of known variable stars in the Galaxy and in the MCs. They also provided the unique opportunity to improve the sampling along the light curves of fundamental (F) and overtone pulsators. In particular, the new multi-band data on classical Cepheids have had a substantial impact not only on the pulsation properties of these variables but also on the estimate of MC distances (Udalski 1998; Groenewegen \\& Oudmaijer 2000, GO00; Groenewegen 2000, G00; Bono, Caputo, \\& Marconi 2001, BCM01). Even though F Cepheids are the most popular standard candles to estimate distances, several theoretical and empirical investigations have been recently focused on Cepheids pulsating in the first (G00; Feuchtinger, Buchler \\& Koll$\\acute{a}$th 2000, FBK00; Baraffe \\& Alibert 2001, BA01) or in the second overtone (BCM01). The main advantage in using overtone pulsators to estimate distances is that the width in temperature of their instability regions is significantly smaller than for the fundamental one. As a matter of fact, current predictions suggest that at $\\log P=0.3$ the width of FO instability strip is 400 K, while for F variables it is 900 K at $\\log P=1$. Therefore, distances based on FO Period-Luminosity (PL) relations are marginally affected by intrinsic spread when compared with F ones. However, we still lack a comprehensive analysis of the uncertainties affecting distance estimates based on optical and NIR PL relations and on the Wesenheit function. A similar approach was already adopted by G00 but in the comparison between theory and observations he was forced to fundamentalize the period, since theoretical predictions for FO in MCs were not available. This gap was filled by BA01 who estimated PL, PLC relations, and Wesenheit function on the basis of linear, convective models. Although FO variables present several undisputable advantages, complete and accurate samples are only available for MC Cepheids. Moreover, the detection of FOs in external galaxies is more difficult than for F Cepheids, since they are fainter and the luminosity amplitudes are smaller.  The main aim of this Letter is to give a detailed analysis of pros and cons of the PL relations currently adopted to estimate the distances, according to full amplitude, nonlinear, convective models of Galactic and MC Cepheids. We are also interested in testing whether current models account for the observed upper limit in the FO period distribution. ",
        "conclusions": "Figure~3 shows the comparison between theory and observations for LMC FO variables. Note that to investigate in detail the accuracy of current models we plotted the blue and the red edge of the instability region instead of mean relations. Data plotted in this figure do suggest that, within current uncertainties, predicted edges agree quite well with empirical data. The agreement found in the $M_I-\\log P$ plane suggests that the discrepancy in the distance moduli based on the $PL_I$ relation is due to a mild overestimate of individual reddening corrections. The dashed lines display the mean relations based on linear models provided by BA01. The comparison discloses that these relations marginally account for empirical data, since they are located close to blue edges of the instability strip. The difference between the two sets of pulsation models is mainly due to fact that BA01 adopted a canonical ML relation. Note that current FO edges do not show, in agreement with empirical data, a change in the slope when moving from 3.25 to 4 $M_\\odot$ (Alibert et al. 1999).  The discrepancies we found are somehow at odds with the results obtained by BA01 concerning the difference between the slopes of their $W$ functions and the slopes predicted by Caputo, Marconi, Musella (2000) for F pulsators.  The main difference is that BA01 were forced to extrapolate the $W$ function given by Caputo et al. (2000). In fact, these relations do rely on F models with periods ranging from $\\log P = 0.50$ to $\\log P=1.86$ for Z=0.004 and from $\\log P=0.51$ to $\\log P=1.93$ for Z=0.008. However, both theoretical predictions (BCM01) and empirical data (Bauer et al. 1999; G00) support the evidence that the slope of the F edges changes when moving from higher to lower luminosities. This means that the extrapolation of PL and $W$ relations toward shorter periods is risky, and therefore we did not perform any selection among unstable F and FO models.   The empirical lower limit in the period distribution of F and overtone Cepheids is a key observable to constrain the accuracy of pulsational and evolutionary models, since it supplies tight constraints on the minimum mass whose blue loop crosses the instability strip, as well as on the topology of the strip (Alcock et al. 1999; Bono et al. 2000; Beaulieu et al. 2001). On the other hand, the upper limit in the period distribution of FO Cepheids is a robust observable to constrain the plausibility of pulsation models (Bono et al. 1999; FBK00). FOs located close to the {\\em intersection point}, i.e. the region of the instability strip where the blue and the red edge of FOs intersect, can be adopted to constrain the luminosity above which Cepheids only pulsate in the F mode. The longest predicted FO period is the aftermath of the topology of the instability strip, and in turn of the physical assumptions adopted to construct the pulsation models. According to empirical evidence based on the OGLE database the longest FO periods range from $\\log P\\approx 0.65$ to 0.77 for the SMC and from $\\log P\\approx0.75$ to 0.80 for the LMC. Empirical estimates based on different Cepheid samples give quite similar upper limits for FOs (Beaulieu \\& Marquette 2000). Note that previous upper limits only bracket 4 (SMC) and 6 (LMC) FO Cepheids respectively. Moreover, the 4 SMC FOs with longer periods present a peculiar position in the Wesenheit plane, and therefore they could have been misclassified. These objects deserve a more detailed analysis to assess whether they are genuine FO pulsators.  To avoid deceptive errors in the comparison between theory and observations we decided to adopt 0.65 and 0.75 as upper limits in the period cut-off.  Current predictions suggest that the cut-off periods for FO Cepheids in MCs are located at $\\log P\\approx0.77$ (SMC) and $\\log P\\approx0.73$ (LMC). The predicted SMC upper limit is $\\approx 0.1$ dex longer than observed, whereas the predicted LMC upper limit agrees quite well with the empirical one. Therefore it is not clear whether this discrepancy is due to a limit of theoretical models. It is noteworthy that pulsation models provided by BA01 predict stable FOs at periods longer than $\\log P\\approx 1.3$ (SMC) and $\\log P\\approx 1.4$ (LMC). These upper limits are substantially longer than observed ones. The reason for this difference is not clear. However, pulsation models based on a similar theoretical framework (linear, nonadiabatic),  and treatment of convective transport, predict similar cut-off periods (Chiosi et al. 1993). To investigate the reasons of the mismatch between predicted and observed SMC FO cut-off periods we compared our predictions with empirical data for Galactic FO Cepheids. Empirical estimates suggest that the longest Galactic FO periods range from $\\log P\\approx 0.7$ to 0.88, but only 3 objects are included in this period range (Kienzle et al. 1999). Current models at solar chemical composition predict a cut-off period, $\\log P\\approx 0.6$, that is $\\approx 0.1$ dex shorter than the observed one. Note that nonlinear, convective models constructed by FBK00 predict an upper limit that is 0.25 dex longer ($\\log P\\approx 0.95$) than observed. Since these models rely on a similar theoretical framework the difference could be due to the different treatment in the turbulent convection model (see \\S 4.1 in FBK00) and/or to the adopted ML relation. Synthetic Color-Magnitude diagrams that account for evolutionary and pulsation properties are mandatory to constrain star formation rates and/or theoretical predictions (Alcock et al. 1999).  The main conclusion we can draw is that current predictions concerning FO periods at the {\\em intersection point} among Galactic and MC Cepheids agree within 0.1 dex with empirical values. A large sample of FO Cepheids in a different stellar system could be crucial to improve the accuracy of the empirical scenario. Accurate radius estimates of Galactic long-period FOs are strongly required to identify their pulsation mode. Different theoretical frameworks predict longer cut-off periods. This means that this observable can be adopted to validate the plausibility of pulsation models.  We thank D. Bersier as a referee for useful suggestions that improved the paper. This work was supported by MIUR/Cofin2000 under the project {\\em Stellar Observables of Cosmological relevance}.   \\pagebreak"
    },
    "0206/astro-ph0206428_arXiv.txt": {
        "abstract": "Using radiative transfer calculations and cosmological simulations of structure formation, we study constraints that can be placed on the nature of the cosmic ultraviolet (UV) background in the redshift interval $2.5\\simlt z \\simlt 5$.  Our approach makes use of observational estimates of the opacities of hydrogen and singly ionised helium in the intergalactic medium during this epoch.  In particular, we model the reionisation of He{\\small II} by sources of hard ultraviolet radiation, i.e. quasars, and infer values for our parameterisation of this population from observational estimates of the opacity of the He{\\small II} Lyman-alpha forest.  Next, we estimate the photoionisation rate of H{\\small I} from these sources and find that their contribution to the ionising background is insufficient to account for the measured opacity of the H{\\small I} Lyman-alpha forest at a redshift $z\\sim 3$.  This motivates us to include a soft, stellar component to the ionising background to boost the hydrogen photoionisation rate, but which has a negligible impact on the He{\\small II} opacity.  In order to simultaneously match observational estimates of the H{\\small I} and He{\\small II} opacities, we find that galaxies and quasars must contribute about equally to the ionising background in H{\\small I} at $z\\simeq 3$.  Moreover, our analysis requires the stellar component to rise for $z > 3$ to compensate for the declining contribution from bright quasars at higher redshift.  This inference is consistent with some observational and theoretical estimates of the evolution of the cosmic star formation rate.  The increasing dominance of the stellar component towards high redshift leads to a progressive softening of the UV background, as suggested by observations of metal line absorption.  In the absence of additional sources of ionising radiation, such as mini-quasars or weak active galactic nuclei, our results, extrapolated to $z > 5$, suggest that hydrogen reionisation at $z\\sim 6$ mostly likely occurred through the action of stellar radiation. ",
        "introduction": "Determining the relative contributions of different sources to the cosmic ultraviolet background is essential for understanding the evolution of the intergalactic medium (IGM).  In particular, this metagalactic radiation field is believed to have reionised hydrogen at $z\\sim 6$ (e.g. Becker et al. 2001) and helium slightly later, although the epoch of helium reionisation has yet to be determined observationally (for a discussion, see e.g. Sokasian et al. 2002). Evidence from measured temperature changes, optical depth variations, and evolution in the relative abundances of metal line absorbers strongly suggests that most intergalactic helium became fully ionised at redshifts close to $\\sim 3.2$ (e.g.  Davidsen et al. 1996, Jakobsen et al. 1994, Kriss et al. 2001, Reimers et al. 1997, Songaila 1998, Ricotti, Gnedin \\& Shull 1999, Theuns et al. 2002a, Theuns et al. 2002b, Bernardi et al. 2002 and references therein).  An important probe of the physical state of the intergalactic medium is provided by bright objects at great distances, such as quasars. For example, it is now believed that absorption by diffuse, cosmologically distributed gas is responsible for the hydrogen Lyman alpha forest (e.g. Cen et al. 1994; Zhang et al. 1995; Hernquist et al.  1996).  Similarly, Ly$\\alpha$ absorption by He {\\small II} along a line of sight to a distant quasar probes gas in the intervening IGM at even lower overdensities (Croft et al. 1997), characteristic of much of the baryonic matter in the Universe (e.g. Dav\\'e et al. 2001; Croft et al. 2001).  At a given redshift, the number and strengths of these spectral features is sensitive to the local density of absorbing atoms, which in turn depends on the gas density, cosmological parameters, and the intensity of the ionising background. In fact, much of the interpretation of spectroscopic observations of high redshift quasars relies strongly on this simple picture of the Lyman alpha forest.  Specifically, given a model for the formation of large-scale structure, the number of lines detected in the Lyman-$\\alpha$ forest as a function of redshift directly constrains the evolution and spectral properties of the radiation field.  In a recent study, Kim, Cristiani, \\& D'Odorico (2001) showed that the number of lines per unit redshift, $dN/dz$, with column densities in the interval $N_{\\text{H {\\tiny I}}}=10^{13.64 - 16}$ decreases continuously from $z\\sim 4$ to $z\\sim 1.5$ according to $dN/dz\\propto (1+z)^{2.19\\pm0.27}$.  Combined with the results of Weymann et al. (1998), who find a much flatter distribution for $dN/dz\\propto (1+z)^{0.16\\pm0.16}$ at $z<1$ , it appears that the line number density of the Ly$\\alpha$ forest is well described by a double power-law with a break at $z\\sim 1$.  These results suggest that the evolution of the forest above $z>1.5$ is governed mainly by Hubble expansion and that there is little change in the ionising background until the break occurring at $z\\sim 1$.  The location of the observed break, however, is inconsistent with theoretical predictions derived from numerical simulations.  In particular, studies of the Lyman-$\\alpha$ forest carried out by Dav\\'e et al. (1999) and Machacek et al. (2000) predict a break in the double power-law occurring near $z\\sim 1.8$.  While these simulations have provided a successful general description of the evolution of the Lyman-$\\alpha$ forest, their apparent inability to match the location of the break indicates that the underlying assumptions regarding the form of the UV background may be incorrect. More specifically, these simulations assume a quasar (QSO) type source population mainly responsible for producing the radiation field. Since the emissivity of quasars is known to fall off steeply below $z \\sim 2$, so would their contribution to the UV background, thereby producing a break in $dN/dz$ around this redshift.  One way to reconcile the inconsistency between the simulations and observations is to appeal to other types of sources to maintain the intensity of the UV background at a relatively high level until $z\\sim 1$.  Recently, Bianchi et al. (2001) have explored the possibility that galaxies might provide this additional contribution to the radiation field.  In particular, they derived the H {\\small I} ionising background resulting from the integrated contribution of quasars and galaxies, taking into account the opacity of the intervening IGM.  The quasar emissivity was derived from fits to an empirical luminosity function, while a stellar population synthesis model and a cosmic star-formation history from UV observations were used to estimate the galaxy emissivity.  They found that the break at $z\\sim 1$ implied by the Kim et al. (2001) analysis can be understood if the contribution from galaxies is comparable to or larger than that of quasars.  This is consistent with other determinations of the galactic component of the background (Giallongo, Fontana, \\& Madau 1997; Devriendt et al. 1998; Shull et al. 1999; Steidel, Pettini, \\& Adelberger 2001).  A significant contribution to the radiation field from galaxies would imply a considerable softening of its spectrum compared to previous models which included only quasars as the dominant source of the ionising metagalactic flux.  In this paper, we shift the focus to higher redshifts to see whether including an additional component from galaxies together with a realistic quasar model is capable of producing the required ionising intensity to match observations of H {\\small I} photoionisation rates in the redshift range $2.5<z<5$.  Our method involves combining numerical and empirical results on the reionisation of singly ionised helium (He {\\small II}) by quasars as a tool for estimating the additional contribution from galaxies to the background which would be necessary to simultaneously match results derived from the H {\\small I} Ly$\\alpha$ forest.  In particular, we select a quasar model based on the study conducted in our earlier paper (Sokasian, Abel \\& Hernquist 2002; hereafter SAH), where we applied a numerical method to study the 3D reionisation of He {\\small II} by quasars. The adopted model is then used to calculate the contribution to the H {\\small I} background from quasars only.  We can then estimate the required contribution from galaxies to match the H {\\small I} photoionisation rates measured by Rauch et al. (1997) in the redshift range $2.5<z<5$.  Consequently, this approach allows us to make estimates of the amplitude and evolution of this component.  This paper is organised as follows.  In Sections 2, 3, and 4 we describe our approach for calculating the emissivities of quasars, galaxies, and recombination radiation from the IGM, respectively.  The procedure for determining the ionising background for hydrogen given these emissivities is presented in Section 5.  Results of our analysis are discussed in Section 6 and conclusions are given in Section 7. ",
        "conclusions": "In this paper, we have utilised observations in both H {\\small I} and He {\\small II} to estimate the contributions to the UV background from quasars and galaxies.  The fact that only quasars are capable of producing radiation hard enough to ionise He {\\small II} has allowed us to select a particular quasar model based solely on the He {\\small II} opacity measurements, independent of the galactic model. Including measurements of the photoionisation rate in H {\\small I} between $2.5\\simlt z\\simlt 4.5$ then enables us to study predictions based on two hotly debated models for the galactic contribution.  We find that a quasar model with less emission than the widely quoted ``standard'' model in conjunction with a galactic model with a slightly rising SFR that contributes a comparable amount of H {\\small I} ionising radiation at $z\\simeq3$, is necessary to achieve good agreement with all relevant observations.  Such a composite model makes it much easier to understand how the intergalactic medium can remain highly ionised at redshifts $z>4$ where the contribution from bright quasars falls off significantly.  It can also explain the apparent progressive softening of the UV background at $z>3$ as suggested by metal absorption line observations (Savaglio et al. 1997; Songaila 1998).  The particular choice for the galactic component in the above model is further bolstered by the fact that the galactic model appears to match observations for the emissivity of LBGs between $2.5\\simlt z\\simlt 4.5$ as it appears as if the adopted galactic model M2 offers good agreement with the observations if the typical value for $f[1500]/f[900]$ lies somewhere between the values measured by Steidel et al. (2001) and Giallongo et al. (2002).  Moreover, the rise in the SFR beyond $z\\sim 3$ is in accord with some observations and theoretical predictions (for a discussion, see e.g. Springel \\& Hernquist 2002a).  We have also shown that there exists a degenerate class of quasar models which are equally successful at matching the He {\\small II} observations while producing a large dispersion in their H {\\small I} contributions.  However, the particular QSO model adopted in this paper has the interesting property of being characterised by plausible values for $L_{\\rm min}$ and $\\alpha_s$ while naturally requiring an extra galactic component that seems to be consistent with observations of both its amplitude and shape.  While the scenario presented in this paper appears promising, we must emphasise that it is based on relatively sparse data.  Future observations gathered with the Sloan Digital Sky Survey (SDSS) will allow us to reduce the uncertainties associated with quasar models at high redshifts.  Coupled with future measurements of the proximity effect and the evolution of the intensity ratio of metal lines, it should soon be possible to place even tighter constraints on the relative contributions from quasars and galaxies to the UV background at $z\\sim 2.5 - 5$."
    },
    "0206/hep-ph0206217_arXiv.txt": {
        "abstract": "The energy spectra of ultra high energy cosmic rays reported by the AGASA, Fly's Eye, Haverah Park, HiRes, and Yakutsk experiments are all shown to be in agreement with each other for energies below $10^{20}$~eV (after small adjustments, within the known uncertainties, of the absolute energy scales). The data from HiRes, Fly's Eye, and Yakutsk are consistent with the expected flux suppression above $5\\times 10^{19}$~eV due to interactions of cosmic rays with the cosmic microwave background, the Greisen-Zatsepin-Kuzmin (GZK) suppression, and are inconsistent with a smooth extrapolation of the observed cosmic ray energy spectrum to energies $> 5\\times 10^{19}$ eV. AGASA data show an excess of events above $10^{20}$~eV, compared to the predicted GZK suppression and to the flux measured by the other experiments. ",
        "introduction": "\\label{sec:introduction}  We analyze the observed spectrum of ultra-high energy cosmic rays. We find two main results: (i) The energy spectra reported by the AGASA, Fly's Eye, Haverah Park, HiRes and Yakutsk experiments are all in good agreement for energies below $10^{20}$ eV, and (ii) All the data are consistent with a GZK suppression except for the AGASA points above $10^{20}$ eV. Our principal conclusion from these two results is that standard physics, including the GZK suppression, is sufficient to explain all of the existing data on UHE cosmic rays.  For any theoretical model in which the GZK suppression is present, the assumed intrinsic spectrum produced by the UHE cosmic-ray sources influences the energy spectrum predicted by the model. Our conclusion that the data are consistent with a GZK suppression implies that the observed spectrum is consistent with model predictions for a plausible intrinsic energy spectrum. In particular, we show that the observed spectrum is consistent with that expected for a GZK suppression of the flux produced by a simple cosmological distribution of sources, each source producing high energy protons with a spectrum $dN/dE_p\\propto E_p^{-2}$ characteristic for collisionless shock acceleration.  Before entering into any details, we will summarize and compare in this introduction the data that are available from different collaborations that measure the spectrum of ultra high energy cosmic rays.   \\subsection{Summary of available data} \\label{subsec:availabledata}  Figure~\\ref{fig:beforeafter} is a \"Before-After\" figure of the currently available data on the highest energy cosmic rays (energies $> 10^{18}$ eV). In Figure~\\ref{fig:beforeafter}a (the \"Before\" version of the figure), the data are plotted, together with their flux error bars, as they have been published by the five experimental collaborations: AGASA~\\cite{agasa}, Fly's Eye~\\cite{fly},  Haverah Park \\cite{HP}, HiRes~\\cite{HiRes}, and Yakutsk \\cite{yakutsk}. The Haverah Park data have recently been re-analyzed using modern numerical simulations of air-shower development~\\cite{HP}. The reanalysis resulted in significant changes of inferred cosmic-ray energies compared to previously published results (\\cite{Watson91} and references quoted therein). The data points for the  Haverah Park measurements that are shown in Fig.~\\ref{fig:beforeafter} are based on this improved analysis, which is available only at energies $<10^{19}$~eV\\footnote{A single flux point at $\\sim 7\\times10^{19}$~eV is shown in Figure~\\ref{fig:beforeafter}a, but this point is based on a preliminary analysis of 4 events that are chosen by different cuts than those applied for the lower energy data. The energy uncertainty for the point at $\\sim 7\\times10^{19}$~eV is significantly larger than the estimated uncertainties for lower energy points~\\cite{HP}. Therefore, the point at  $\\sim 7\\times10^{19}$~eV is shown in Figure~\\ref{fig:beforeafter}a only for completeness; it is not used elsewhere in our analysis because the Haverah Park collaboration has described this point as preliminary.}.  The most striking feature of Figure~\\ref{fig:beforeafter}a is that the experimental results differ greatly among themselves (by factors $\\sim2$) even in the region $10^{18}$~eV $< E < $ $2\\times10^{19}$~eV, where the quoted error bars from each experiment are very small. In addition, the higher AGASA flux reported above $2\\times10^{20}$~eV stands out above the scatter in the different experimental measurements.  Figure~\\ref{fig:beforeafter}b (the \"After\" version of our \"Before-After\" figure) shows a dramatically different representation of the available data. With small adjustments in the absolute energy scales, all of the measured fluxes are seen to be in agreement at energies below $10^{20}$ eV. In constructing Figure~\\ref{fig:beforeafter}b, we have adjusted the absolute energy calibrations within the error bars published by the experimental collaborations. We chose the shifts so as to bring the different measured fluxes into agreement at $10^{19}$~eV. The energy shifts can be accomplished in five equivalent ways, depending upon which one of the five energy scales is unaltered. For Figure~\\ref{fig:beforeafter}b, the Fly's Eye energy scale was unaltered and  we adjusted the AGASA energy scale by -11\\%, Haverah Park by +15\\%, HiRes by +7.5\\%, and Yakutsk by -19\\%. All shifts are well within the published systematic errors.  Figure~\\ref{fig:beforeafter} illustrates visually our two main points. First, all of the currently available data on high energy cosmic rays are in agreement within their quoted errors for energies between $2\\times 10^{18}$~eV and $10^{20}$~eV. Second, three of the four data sets available above $10^{19}$~eV, HiRes, Fly's Eye, and Yakutsk, all show evidence for a turnover of the energy spectrum for energies above $5\\times10^{19}$~eV. This turnover, we shall show later, is highly significant statistically and is consistent with what one would expect from a simple model that includes the GZK effect.  Above $10^{20}$~eV,  the reported AGASA fluxes are higher than the fluxes measured in other experiments. It is these high AGASA fluxes alone that have led to the widespread impression that measurements of ultra-high energy (UHE) cosmic rays (energies $>10^{19}$~eV) do not show evidence for a GZK effect.   \\subsection{What does it all mean?} \\label{subsec:whatmean}  What can one make of the results shown in the Before-After Figure~\\ref{fig:beforeafter}?  There are two simple possibilities. First, the excellent agreement shown in Figure~\\ref{fig:beforeafter}b among the different experiments could be accidental. According to this interpretation, the small adjustments made in the energy scales are not physically motivated and the real situation is somehow much more complicated. It is just a fluke that all of the adjusted energy spectra line up together so well below $10^{20}$~eV. This interpretation is certainly possible. In the present paper, however, we shall choose a different interpretation of Figure~\\ref{fig:beforeafter}b. We shall suppose that the excellent agreement of the adjusted energy spectra reveals a good approximation to the true shape of the UHE cosmic ray energy spectra. We shall now explore the consequences of this assumption.  We stress that the distinction between the two possibilities for interpreting Figure~\\ref{fig:beforeafter}b can only be settled by a new generation of precise and high statistics measurements of the UHE cosmic ray spectrum. Fortunately, the Auger experiment, currently under construction~\\cite{auger}, is expected to provide the necessary precision and statistics. The Telescope Array experiment~\\cite{TA}, currently under planning, may also provide similar precision and statistics.  We first describe in Section~\\ref{sec:model} the model we use and then in Section~\\ref{sec:comparison} we compare the model predictions with observations of the UHE cosmic ray energy spectrum. We summarize our main conclusions in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  Our most important conclusion is that exotic new physics is not required to account for the observed events with energies in excess of $10^{20}$ eV, except for the AGASA data. Table~\\ref{tab:predictedobserved} shows that there is already a strong suggestion, $> 5\\sigma$ ($>3.7\\sigma$, depending upon the extrapolated energy spectrum) in the Fly's Eye, HiRes, and Yakutsk observations that the expected GZK suppression has been observed (see also Fig.~\\ref{fig:beforeafter} and Fig.~\\ref{fig:spectrumvsmodel}).  Precision measurements from $10^{18}$~eV to $5\\times 10^{19}$~eV are essential for testing models of UHE cosmic rays, although they are less dramatic than measurements above $10^{20}$~eV. At energies $> 10^{20}$ eV, the predicted number, $N$, of events in conventional models is uncertain due to the unknown clustering scale, $r_0$, of the sources, $\\sigma(N_{\\rm predicted})/N_{\\rm predicted} = 0.9(r_0/{10~\\rm Mpc})^{0.9}$ ~\\cite{clustering}. Paradoxically, we may need to study carefully cosmic rays with energies below the GZK suppression in order to understand better the origin of the cosmic rays beyond the suppression."
    },
    "0206/astro-ph0206178_arXiv.txt": {
        "abstract": "The Geneva group has reported two Saturn--mass planets orbiting HD 83443 (K0V) with periods of 2.98 and 29.8 d.  The two planets have raised interest in their dynamics because of the possible 10:1 orbital resonance and the strong gravitational interactions.  We report precise Doppler measurements of HD 83443 obtained with the Keck/HIRES and the AAT/UCLES spectrometers.  These measurements strongly confirm the inner planet with period of 2.985 d, with orbital parameters in very good agreement with those of the Geneva group. However these Doppler measurements show no evidence of the outer planet, at thresholds of 1/4 (3 \\ms) of the reported velocity amplitude of 13.8 \\ms.  Thus, the existence of the outer planet is in question.  Indeed, the current Doppler measurements reveal no evidence of any second planet with periods less than a year. ",
        "introduction": "\\label{intro}  Several multiple planet systems have been reported, including the triple planet system around Upsilon Andromedae (Butler et al. 1999) and double planet systems around GJ876 (Marcy et al. 2001b), HD 83443 (Mayor et al. 2000), HD 168443 (Marcy et al. 2001a, Udry et al. 2002) and 47 UMa (Fischer et al. 2002).  Double-planet systems have also been reported in a press release (ESO Press Release 07/01, 2001) for HD 82943 and HD 74156.  These multiple planet systems contain planets reported to range from a saturn mass to nearly 10 jupiter mass, all orbiting within 4 AU.  Interactions between the planets in some of these systems, notably Gliese 876, are measurable on a time scale of a few years (Lissauer \\& Rivera 2001; Laughlin \\& Chambers 2001; Rivera \\& Lissauer 2001). Doppler measurements can reveal the ongoing gravitational perturbations and constrain both the planet masses and orbital inclinations.  The interactions and orbital resonances, both mean--motion and secular, provide clues about the dynamical history of the systems (Snellgrove et al. 2001; Lee and Peale 2002a; Chiang et al. 2002).  A most extraordinary double planet system was reported for HD 83443 (Mayor et al. 2000).  Their Doppler measurements made with the CORALIE spectrometer indicate the existence of two saturn--mass planets that both reside within 0.2 AU.  The inner planet has an orbital period of 2.985 d, an eccentricity of 0.079 ($\\pm$0.033), a minimum (\\msini) mass of 0.34 \\mjup, and an orbital distance of 0.038 AU.  The orbital period is the shortest known for extrasolar planets.  The non--zero eccentricity of this inner planet is notable as planets with periods less than 5 d suffer tidal circularization (Wu \\& Goldreich 2002). With the exception of the saturn--mass planet around HD 46375 (Marcy et al. 2000), all 15 of the previously discovered ``51 Peg--like'' planets have spectral types of G5 or earlier.  Mayor et al. (2000) report a remarkable second planet around HD 83443. It has an orbital period of 29.83 ($\\pm$0.18) d, an eccentricity of 0.42, a minimum (\\msini) mass of 0.16 \\mjup, and an orbital distance of 0.17 AU.  This outer planet induces a velocity semiamplitude in the star of $K$ = 13.8 $\\pm$ 1 \\ms, rendering it a 14-$\\sigma$ detection. This planet has the smallest \\msini yet reported, and is only the third reported planet with a semiamplitude smaller than 15 \\ms (cf. Marcy et al. 2000; Fischer et al. 2002).  Both of the planets were indicated by Doppler measurements obtained with the 1.2-m Leonhard Euler telescope at the ESO La Silla Observatory, which feeds the CORALIE spectrometer (Queloz et al. 2000a).  Wavelength calibration is achieved by coupling the telescope and thorium lamp to the spectrometer with a double scrambled fiber.   The quoted instrumental precision is now 2 \\ms (Udry et al. 2002).   As the two planets orbiting HD 83443 are crowded within 0.2 AU, the system is dynamically active.  Calculations by J.Laskar and W.Benz (reported in Mayor et al. 2000) and by Wu \\& Goldreich (2002) and Lee \\& Peale (2002b) suggest the occurrence of significant gravitational interactions between the two planets.  The tidal circularization time scale for the inner planet to HD 83443 is estimated to be $3 \\times 10^{8}$ yr (Wu \\& Goldreich 2002), while the star is estimated to have an age of 6.5 Gyr. In this context, the non--zero eccentricity of the inner planet and the apside alignment of the two orbits are understood to be due to secular interactions between the two planets and tidal interactions with the star (Mayor et al. 2000; Wu \\& Goldreich 2002).  These in turn constrain the orbital inclination of this system, and the radius of the inner planet (Wu \\& Goldreich 2002).  The dynamical evolution that led to the system may involve migration and resonances (Lee \\& Peale 2002b).  Section 2 of this paper will describe new Doppler measurements of HD 83443 made from the Keck and AAT telescopes.  Section 3 describes a search for the two planets, notably the interesting outer planet.  Our failure to detect the outer planet is discussed in Section 4. ",
        "conclusions": "Precision Doppler observations made with the 10--m Keck and the 3.9--m AAT strongly confirm the existence of the inner planet orbiting HD 83443, and indicate the orbital parameters are in very good agreement with those reported by Mayor et al. (2000). The present orbital parameters differ only marginally in that the orbit of the inner planet is circular within measurement uncertainty for the Keck and AAT data, similar to the other known close--in planets.  But our Doppler measurements did not detect the 29.8 d outer planet, despite the clear ability to do so. The present measurements impose a limit on any such velocity periodicity at a level of no more than 3 \\ms, well below the reported velocity amplitude of 13.8 \\ms.  Orbital periods within 2 days of 29.8 d would have been detected.  The supposed velocity amplitude of 13.8 \\ms is four times larger than the uncertainties in our velocity measurements rendering the outer planet immediately detectable.  Various tests suggest quantitatively that the velocities should have revealed the outer planet.  The velocities from the Keck and AAT telescopes could have independently detected the outer planet, but neither data set revealed it.  We considered various possible reasons that we failed to detect the outer planet.  One possibility is that some interactive resonance between the two planets causes the reflex velocity of the star to mimic insidiously a single Keplerian orbit. That is, perhaps the 10:1 ratio of orbital periods, along with gravitational interactions, yields a final reflex velocity that traces a single Keplerian velocity curve.  If so, we might be fooled into fitting the velocities with such a simple model.  We find this possibility unlikely.  As shown by W. Benz (Mayor et al. 2000) and by Wu \\& Goldreich (2002), the gravitational interactions yield temporal evolution of the orbits on a time scale of $\\sim$1000 yr rather than a few years.  Thus we expect the outer planet, if it exists, to remain in a coherent orbit during the few year duration of the present observations.  Moreover, the 10:1 ratio of the two periods do not constitute a powerful Fourier harmonic from which a single Keplerian may be constructed (as is the case with a 2:1 ratio of periods).  We remain puzzled by the discrepancy between the reported CORALIE results and the velocities we have obtained with Keck/HIRES and AAT/UCLES.  Of the 75 extrasolar planet candidates (\\msini $<$ 13 \\mjup) announced from precision Doppler surveys\\footnote{$http://exoplanets.org/almanacframe.html$}, a total of 57 have been published in refereed journals. These planets are listed in Table 3, which also notes the telescope from which the data originates, and whether the actual Doppler velocities are publicly available. Refereed precision velocity confirmations are also included. Substellar candidates found by other techniques such as astrometry and low precision Doppler velocities are not included. An additional 4 planet candidates have been announced in conference proceedings\\footnote{HD 6434, HD 19994, HD 121504, HD 190228} (Queloz et al. 2000b; Sivan et al. 2000).  Doppler velocity measurements are not available for candidates that have only been published in conference proceedings.  An additional 12 claimed Doppler planets, all of which were announced more than 1 year ago, have not been submitted to either a conference proceeding or a refereed journal.  While the discovery of extrasolar planets has become seemingly commonplace over the past 6 years, we still consider the detection of planets orbiting other stars as extraordinary, and as such worthy of the dictum, ``Extraordinary claims require extraordinary evidence.''  Publishing discovery data in a refereed journal remains a crucial part of the process, though this is not in itself sufficient to establish the credibility of a planet claim.  It remains extremely likely that at least a handful of the reported planets do not in fact exist. Multiple confirmation both by independent precision Doppler teams and by completely independent techniques remain the only means by which to ensure the veracity of extrasolar planet claims."
    },
    "0206/astro-ph0206452_arXiv.txt": {
        "abstract": "The observed properties of novae before and after eruption are discussed. The distribution of orbital periods of novae shows a concentration near 3.2 h, which resembles that of magnetic cataclysmic variables, and there is some evidence that many of the novae themselves are magnetic near that orbital period. Desynchronisation of polars by nova eruptions can lead to an estimate  ($\\sim 2 \\times 10^3$ y) for the time between eruptions for the strongly magnetic systems; this is much shorter than that found from other methods. The similarity of pre- and post-nova luminosities, at high rates of mass transfer, is ascribed to irradiation of the secondary producing a self-sustained high $\\dot{M}$ state. This slows cooling of the white dwarf after eruption, delays the onset of full scale dwarf nova outbursts in most systems, and delays any descent into a hibernation state of low rate of mass transfer. ",
        "introduction": "This conference is mostly about the high luminosity state, and the transitions into and out of it, but full understanding of the nova process must include the nature of the low luminosity white dwarf and its accretion environment, in which it spends most of its life. ",
        "conclusions": ""
    },
    "0206/astro-ph0206208_arXiv.txt": {
        "abstract": "The HI disk of the blue compact dwarf (BCD) galaxy NGC 2915 extends to 22 optical scalelengths and shows spiral arms reaching far beyond the optical component. None of the previous theories for spiral structure provide likely explanations for these very extended spiral arms.  Our numerical simulations first demonstrate that such large spiral arms can form in an extended  gas disk embedded in a massive triaxial dark matter halo with slow figure rotation, through the strong gravitational torque of the rotating halo.  We then show that the detailed morphological properties of the developed spirals and rings depend strongly on the pattern speed of the figure rotation, the shape of the triaxial halo, and the inclination of the disk with respect to the plane including the triaxial halo's long and middle axes.  These results strongly suggest that the dark matter halo of NGC 2915 is triaxial and has figure rotation.  Based on these results, we also suggest that dynamical effects of triaxial halos with figure rotation are important in various aspect of galaxy formation and evolution, such as formation of polar ring galaxies, excitation of non-axisymmetric structures in low surface-brightness galaxies, and gas fueling to the central starburst regions of BCDs. ",
        "introduction": "To understand the  distribution and the nature of dark matter in galaxies and dynamical effects of dark matter halos on galaxies has been a longstanding and remarkable problem in galactic astronomy (e.g., Trimble 1987; Ashman 1992; Salucci \\& Persic 1997). Detailed analysis of rotation curves in variously different galaxies have played a major role in revealing the radial distributions of dark matter halos in galaxies (Salucci \\& Persic 1997 for a review). Recently high-resolution rotation curve studies for low surface brightness galaxies and dwarfs have extensively discussed whether or not the radial density profiles of dark matter halos predicted from cold dark matter (CDM) models  (Navarro, Frenk, \\& White 1996) are consistent with  the observationally inferred profiles (e.g., de Blok et al. 2001; van den Bosch \\& Swaters 2001).   Several attempts have been so far made to reveal the {\\it shapes} (e.g., the degree of oblateness or triaxiality) of dark matter halos in galaxies (e.g., Trimble 1987; Ashman 1991). Following the early attempts to use the kinematics of polar rings for deriving the three dimensional mass distributions of galaxies (Schweizer et al. 1983;  Whitmore et al. 1987), Sackett \\&  Sparke (1990) tried to give strong constraints on the shape of a dark matter halo in polar ring galaxy NGC 4650A by investigating both a rotation curve of a polar ring component and that of a planner disk one. They concluded that the best modeled flattening of the dark halo in NGC 4650A is somewhere between E3 and E7. Franx, van Gorkom, \\& de Zeeuw (1994) analyzed both the geometry and the velocity of the  HI gas ring of IC 2006 and found a nonsignificant ellipticity of the gravitational potential ($\\sim$ 0.012 $\\pm$ 0.026).  Furthermore, Olling (1995) found that the thickness of the HI gas disk extending beyond the Holmberg radius in a galaxy is sensitive to the flattening of the dark matter halo and proposed that the high resolution HI studies on  the flaring of the outer gas layers in nearby galaxies enables us to determine the shape of the dark matter halos. By comparing numerical simulations of dynamical evolution of the Sagittarius dwarf with observations on the detailed spatial distribution of the 75 Galactic halo stars, Ibata et al. (2001) suggested that the Galactic dark halo is most likely almost spherical in the Galactocentric distance 16 $<$ $R$ $<$ 60 kpc. Based on the structural  and kinematic properties of the Galactic luminous A-type stars revealed by  Hipparcos data (ESA, 1997), Cr\\'ez\\'e et al. (1998) argued that there are strong lower limits on the scaleheight of any flattened component of the galactic dark halo.   Although the radial density profiles and the shapes of dark matter halos have been extensively discussed by many authors, their {\\it rotational properties} have been less discussed.  Based on the detailed analysis of structure and kinematics of the very extended HI disk around NGC 2915, Bureau et al (1999) first suggested that the observed spiral-like structures in the HI disk can be formed  by a triaxial halo with figure rotation. They furthermore pointed out that the slow pattern speed of the figure rotation inferred from the kinematics and density distribution of the HI gas is consistent with the pattern speeds of rotating triaxial dark halos seen in CDM simulations analysed by Pfitzner (2000).  However, it is unclear whether a triaxial dark halo with figure rotation is really responsible for the observed extended spiral structures in NGC 2915, because of the lack of numerical studies of gas dynamics in the gravitational potentials of triaxial halos with figure rotation.  The purpose of this Letter is to demonstrate how non-axisymmetric structures (e.g. spirals and bars) can be formed in a gas disk well outside the  optical radius of a galaxy embedded in a massive triaxial dark matter halo with figure rotation.  We particularly investigate how  morphological evolution of outer gas disks depends on the structure of the triaxial halos, the  pattern speeds of the figure rotation, and the physical properties of the disks (e.g. inclination of the disks with respect to the halo and gaseous mass and temperature). ",
        "conclusions": "Here are some of the implications of our results. Firstly, some of the morphological properties of polar ring galaxies result from gas response to their triaxial dark halo with figure rotation.  Our models showed that the spatial distributions of gas disks can be strikingly similar to the double-ring structure observed in polar ring galaxies ESO 474 - G26 and NGC 2685 (Schweizer et al. 1983) and to a parallelogram-like structure in  NGC 660 (van Driel et al. 1995).  The simulated structures which appear to be polar ring components in polar ring galaxies are actually not rings but disks with spiral arms and with inner low-density gas regions.  This  implies that {\\it some} polar ring components are not rings but `polar disks' with spiral arms.  NGC 4650A is observed to be such a polar disk with spiral structure (e.g., Arnaboldi et al.  1997), and accordingly might be formed by the torque of its rotating triaxial halo.  Secondly, we recall Davies's (1972) suggestion that the HVC complex C represents an outer extended Galactic spiral arm at relatively high latitude.  The present simulation successfully produced extended outer gas arms in a galaxy embedded in a rotating triaxial halo.  We suggest that high-density regions in the extended gaseous arms excited by the Galactic triaxial halo could be the formation sites of HVCs.  Extensive discussions on the origin of HVCs based on the comparison of our simulations and HIPASS observations will be given in our forthcoming papers (Bekki \\& Freeman 2002).   Thirdly, the morphological properties of low surface brightness (LSB) disk galaxies can be greatly influenced by triaxial dark matter halos with  figure rotation, because these LSBs are observed to be dominated by dark matter (e.g. de Blok \\&  McGaugh 1996).  Our simulations show that even a compact low-mass gas disk can be transformed into spiral arms by the gravitational torque of its rotating triaxial dark matter halo.  The observed diversity in non-axisymmetric structures in LSBs (McGaugh 1992) with possibly large Toomre (1964, 1981) Q and $X$ parameters (i.e. systems in which {\\it spontaneous} spiral formation is highly unlikely) can be due partly to the figure rotation of their triaxial halos.  Our results also imply that even in high surface brightness (HSB) galaxies, {\\it the outer gas} can be strongly affected, if the gas is misaligned with the triaxial halos.  We suggest that the observed  gaseous warps in HSBs can be formed by the halos at the radius where the strong vertical restoring force of {\\it stellar disks} drops rapidly (e.g., Briggs 1990).  Fourthly, the large-scale torques associated with the rotating triaxial halos provide a mechanism for fueling the starburst regions of BCDs by replenishing gas from their outer HI gas reservoirs.   Most dwarf irregular galaxies are observed to have HI gas extending to $\\sim$ twice the Holmberg radius ($R_{\\rm H}$). Some of them have gas out to $4-7 \\times R_{\\rm H}$ (e.g Hunter 1997).  The morphologies of this extended HI gas are diverse, ranging from smooth and quiescent (e.g. NGC 6822; Roberts 1972) to spectacularly complex (e.g. NGC 4449; Bajaja et al. 1994) and with arms or blobs as in DDO 26 (Hunter \\& Wilcots 2002).  Extended HI with spiral-like structures and arcs is seen in spirals like NGC 628 (e.g., Kamphuis \\& Briggs 1992).  Our study strongly suggests that these apparently peculiar structures in the extended outer HI gas of galaxies can tell us something important about the shapes and the rotational properties  of dark matter halos.  Systematic comparison between hydrodynamical simulations such as those presented in this study and high-resolution morphological studies of unusually extended  HI gas of galaxies might well reveal the dependence of the pattern speed of rotating triaxial dark halos on the masses, structures, and environments of galaxies. This could give useful constraints on theories of galaxies formation.  Finally we would welcome high-resolution cosmological simulations which provide better statistics on the incidence of rotating figures of dark matter halos in low-mass dwarf galaxies."
    },
    "0206/astro-ph0206092_arXiv.txt": {
        "abstract": "We present high signal-to-noise integrated spectra of 24 star clusters in the Large Magellanic Cloud (LMC), obtained using the FLAIR spectrograph at the UK Schmidt telescope. The spectra have been placed onto the Lick/IDS system in order to test the calibration of Simple Stellar Population (SSP) models (Maraston \\& Thomas 2000; Kurth, Fritz-von Alvensleben \\& Fricke 1999).  We have compared the SSP-predicted metallicities of the clusters with those from the literature, predominantly taken from the Ca-Triplet spectroscopy of Olszewski \\etal (1991). We find that there is good agreement between the metallicities in the range --2.10 $\\leq$ [Fe/H] $\\leq$ 0. However, the Mg$_2$ index (and to a lesser degree Mg $b$) systematically predict higher metallicities (up to +0.5 dex higher) than $\\langle$Fe$\\rangle$. Among the possible explanations for this are that the LMC clusters possess [$\\alpha$/Fe] $>$ 0. Metallicities are presented for eleven LMC clusters which have no previous measurements.  We compare SSP ages for the clusters, derived from the H$\\beta$, H$\\gamma$ and H$\\delta$ Lick/IDS indices, with the available literature data, and find good agreement for the vast majority. This includes six old globular clusters in our sample, which have ages consistent with their \\HST\\ CMD ages and/or integrated colours. However, two globular clusters, NGC~1754 and NGC~2005, identified as old ($\\sim$ 15 Gyr) on the basis of \\HST\\ CMDs, have H$\\beta$ line-strengths which lead ages which are too young ($\\sim$ 8 and $\\sim$ 6 Gyr respectively). These findings are inconsistent with their CMD-derived values at the 3$\\sigma$ level. Comparison between the horizontal branch morphology and the Balmer line-strengths of these clusters suggests that the presence of blue horizontal branch stars has increased their Balmer indices by up to $\\sim$ 1.0 \\AA.  We conclude that the Lick/IDS indices, used in conjunction with contemporary SSP models, are able to reproduce the ages and metallicities of the LMC clusters reassuringly well. The required extrapolations of the fitting-functions and stellar libraries in the models to younger ages and low metallicities do not lead to serious systematic errors. However, due to the significant contribution of horizontal branch stars to Balmer indices, SSP model ages derived for metal-poor globular clusters are ambiguous without {\\it a priori} knowledge of horizontal branch morphology. ",
        "introduction": "\\label{Introduction}  One of the most powerful tools available to observers of stellar populations is the colour-magnitude diagram (CMD). Whilst there still remain numerous uncertainties in stellar evolution theory (e.g. \\citeANP{Castellani01} 2001), the existence of accurate paralaxes such as those from \\HIPPARCOS, used in conjunction with contemporary model isochrones can now constrain the ages of Galactic globular clusters (GCs) to within $\\sim$ 20\\% (e.g. \\citeANP{Carretta00} 2000).  Furthermore, \\HST\\ has allowed us to obtain detailed information for GCs and field stars in external galaxies such as the LMC (e.g. \\citeANP{Holtzman97} 1997; \\citeANP{Olsen98} 1998; \\citeANP{Johnson99} 1999), the Andromeda galaxy (e.g. \\citeANP{Holland96} 1996; \\citeANP{FusiPecci96} 1996) and even the nearest large elliptical Centaurus A (e.g. \\citeANP{Soria96} 1996; \\citeANP{Harris98} 1998; \\citeANP{GHarris99} 1999; \\citeANP{Marleau00} 2000).  However, such observations are challenging, and beyond several Mpc exceed the capabilities of current instrumentation. Even with \\HST, a combination of crowding and the intrinsic faintness of single stars limits the applicability of such an approach. Therefore, in order to probe the properties of distant stellar systems, we must rely upon studies of integrated light.  Integrated spectroscopy and photometry require comparisons with stellar population models, and are affected by a degeneracy between age and metallicity (\\citeANP{Faber72} 1972; \\citeANP{OConnell76} 1976). Spectroscopic indices have been shown to hold potential, and much work in the past decade has lead to age-metallicity diagnostics for integrated spectra (\\eg Gonz$\\acute{a}$lez~1993; \\nocite{Gonzalez93}\\citeANP{Rose94} 1994; \\citeANP{Worthey94} 1994; \\citeANP{Borges95} 1995; \\citeANP{Idiart95} 1995; \\citeANP{WO97} 1997; \\citeANP{Vaz99} 1999). These methods have subsequently been used by many workers to derive ages and metallicities for galaxies (\\eg \\citeANP{Davies93} 1993; Gonz$\\acute{a}$lez~1993; \\citeANP{Fisher95} 1995; \\citeANP{Harald98} 1998; \\citeANP{Vazdekis01a} 2001) and extragalactic globular clusters (\\eg \\citeANP{Cohen98} 1998; \\citeANP{KisslerPatig98} 1998; \\citeANP{Beasley00} 2000; \\citeANP{Forbes01b} 2001).  However, the reliability of such integrated techniques has yet to be demonstrated: they must be tested against stellar systems with independently derived age and metallicity estimates such as Galactic GCs. Addressing this issue, \\citeANP{Gibson99} (1999) derived a 'spectroscopic' age for 47 Tucanae in the H$\\gamma_{\\rm HR}$ -- Fe4668 and H$\\gamma_{\\rm HR}$ -- CaI$_{\\rm HR}$ planes of the \\citeANP{Worthey94} (1994; henceforth W94) simple stellar population (SSP) models. These authors found that 47 Tuc's integrated spectrum fell below the oldest (17 Gyr) isochrones of the W94 models at the 4$\\sigma$ level, yielding an extrapolated age in excess of 20 Gyr. By comparison, the CMD-derived age of 47 Tuc is 14 $\\pm$ 1 Gyr \\cite{Richer96}. On the other hand, \\citeANP{Maraston00} (2000) derived an age of 15 Gyrs for this cluster using the combination H$\\beta$ and Fe5335, in good agreement with the CMD of \\citeANP{Richer96} (1996). \\citeANP{Vazdekis01} (2001) and \\citeANP{Schiavon02} (2002) have recently addressed these issues and conclude that the inclusion of atomic diffusion and non-solar abundance ratios are important. Moreover, \\citeANP{Schiavon02} (2002) found it necessary to include both AGB stars and adjust the metallicity of 47 Tuc by --0.05 dex to fit their SSP models to the integrated spectrum of this cluster.  Whilst these developments are extremely promising, the full calibration of SSP models has yet to be comprehensively tested. The Galactic GC 47 Tuc represents a single age and single metallicity in the large parameter space of contemporary SSP models. In view of the adjustments employed Schiavon \\etal 2001 in order to reproduce the integrated spectrum of just this cluster, begs the question of how well can these models be applied to more distant, less well-known stellar systems? Furthermore, 47 Tuc is an idealised case of an old, relatively metal-rich stellar system which has a 'red clump' for its horizontal branch (HB). In this case, the strength of the Balmer lines are thought to be primarily a function of the temperature of the main sequence turn-off (and hence age). At lower metallicities\\footnote{but see \\citeANP{Rich97} (1997) for two examples of relatively metal-rich Galactic GCs with blue HBs.} GCs develop blue HBs which are expected to contribute a significant component to Balmer lines (\\eg \\citeANP{Buzzoni89} 1989; \\citeANP{deFreitasPacheco95} 1995; \\citeANP{Lee00a} 2000; \\citeANP{Maraston00} 2000).  On the observational side, obtaining integrated spectra of Galactic GCs may hide other uncertainties. Owing to the large angular size of Galactic GCs on the sky (the half-light radius of 47 Tuc is $\\sim$ 2.8 $\\arcmin$) a spectroscopic aperture must be synthesised by physically scanning a slit across the cluster. In so doing, the observer may unwittingly include foreground stars in the integrated spectrum, in addition to increasing the liklihood of stochastic contributions from bright stars. Clearly, alternative laboratories are desirable to test SSP models.  The Large Magellanic Cloud (LMC) presents an ideal target for such tests. Many of its $\\sim$ 2000 star clusters \\cite{Olszewski96} have been independently age-dated and their metallicities determined, whilst its proximity ($\\sim$ 53 \\kpc) means that acquiring high S/N, integrated spectra of the clusters is relatively straightforward (\\eg \\citeANP{Rabin82} 1982).  In this paper, we present high S/N integrated spectra for 24 star clusters in the LMC, covering a wide range in age (0.5 -- 17 Gyr) and metallicity (--2.1 $\\leq$ [Fe/H] $\\leq$ 0). We have placed the clusters onto the Lick/IDS system and measured their line-strength indices in order to test contemporary stellar population models which use the Lick/IDS fitting-functions (\\citeANP{Gorgas93} 1993; \\citeANP{Wortheyetal94} 1994; \\citeANP{WO97} 1997).  The plan of this paper is as follows: In Section~\\ref{sec:SampleSelection}, we describe our sample selection and the observations performed for this present study. In Section~\\ref{sec:DataReduction} we describe the reduction steps required for our fibre spectra. We discuss the spectroscopic system and the stellar population models we use in this study in Section~\\ref{sec:TheSpectroscopicSystem}. In Section~\\ref{sec:TheAgesandMetallicitiesoftheLMCClusters}, we derive ages and metallicities for the LMC clusters using the SSP models, which we then compare to literature values. Finally, we present our conclusions and a summary in Section~\\ref{sec:SummaryandConclusions}. ",
        "conclusions": "\\label{sec:SummaryandConclusions}  We have obtained high S/N integrated spectra for 24 star clusters in the LMC, and derived age and metallicity estimates for these clusters using a combination of the Lick/IDS indices and the SSP models of \\citeANP{Maraston00} (2000) and \\citeANP{Kurth99} (1999). To test the SSP models, we have compiled a list of metallicity and age determinations from the literature. Metallicities were taken largely from the Ca-Triplet spectroscopy of \\citeANP{Olszewski91} (1991). The age estimates are somewhat more inhomogeneous in nature, deriving from a number of methods; the location of the main-sequence turn-off in CMDs, the extent of the AGB calibration of \\citeANP{Mould82} (1982) or integrated colours. Comparing these independently derived quantities we find:   \\begin{itemize}  \\item the SSP-derived metallicities, obtained using both the Mg$_2$ and $\\langle$Fe$\\rangle$ Lick/IDS indices show generally good agreement with the literature values. However, the Mg$_2$ index (and to a lesser degree Mg $b$) predicts metallicities which are systematically higher than those from $\\langle$Fe$\\rangle$, by up to +0.5 dex for the highest metallicity clusters. Amongst the possible explanations for this difference is the existence of [$\\alpha$/Fe] $>$ 0 in the clusters. We publish metallicities for 11 LMC star clusters with no previous measurements, which are accurate to $\\sim$ 0.2 dex.  \\item for the majority of the LMC clusters, the SSP-models predict ages from the H$\\beta$, H$\\gamma_{\\rm F}$ and H$\\delta_{\\rm F}$ indices which are consistent with the literature values. However, age estimates of the old LMC globular clusters are often ambiguous. The oldest isochrones of the SSP models overlap younger isochrones due to the modelling of mass-loss on the RGB. {\\it Assuming} old ages in interpreting these data, six clusters, NGC~1786, NGC~1835, NGC~1898, NGC~1916, NGC~1939 and NGC~2019, have SSP-derived ages in all three measurable Balmer indices which are consistent with their ages derived from CMDS or integrated colours.  \\item two clusters, namely NGC~1754 and NGC~2005, have extremely strong Balmer lines, which leads to the SSP model ages which are too young ($\\sim$ 8 and 6 Gyr respectively). Comparison between the horizontal branch morphology and the Balmer lines for five of the GCs in our sample suggests that blue HBs are likely contributing up to $\\sim$ 1.0 \\AA\\ to the H$\\beta$ index in these clusters.   \\end{itemize}  We conclude that the SSP models considered in this study are able to satisfactorily predict the ages and metallicities for the vast majority of LMC star clusters from integrated spectroscopic indices. This remains true despite the rather inhomogenous nature of the literature age determinations. However, estimating the ages of the old, low-metallicity LMC GCs is severely complicated by the strong contribution of horizontal branch/post-horizontal branch stars to integrated indices. We conclude that at old ages and low metallicities ([Fe/H] $<$ --1.0), Balmer lines (\\eg H$\\beta$, H$\\gamma$, H$\\delta$) are not useful age indicators without {\\it a priori} knowledge of horizontal branch morphology."
    },
    "0206/astro-ph0206217_arXiv.txt": {
        "abstract": "We show how the morphological analysis of the maps of the secondary CMB anisotropies can detect an extended period of ``smoldering'' reionization, during which the universe remains partially ionized. Neither radio observations of the redshifted 21cm line nor IR observations of the redshifted Lyman-alpha forest will be able to detect such a period. The most sensitive to this kind of non-gaussianity parameters are the number of regions in the excursion set, $N_{cl}$, the perimeter of the excursion set, $P_g$, and the genus (i.e.\\ '1 - number of holes') of the largest (by area) region. For example, if the universe reionized fully at $z=6$, but maintained about $1/3$ mean ionized fraction since $z=20$,then a $2$ arcmin map with $500^2$ pixel resolution and a signal-to-noise ratio $S/N=1/2$ allows to detect the non-gaussianity due to reionization with better than 99\\% confidence level. ",
        "introduction": "Recent observations of high redshift quasars \\cite{bec01,djo01} offer a unique probe of the physical conditions in the intergalactic medium (IGM) shortly after the epoch of reionization. Similar future observations will increase the amount of observational data multi-fold and will provide critical constraints on the theories of galaxy formation.  However, if we want to go beyond the epoch of reionization and to study the pre-overlap stage during which ionized \\hii\\ regions expand into still neutral low density gas, we need to use different bands of electromagnetic spectrum. Next Generation Space Telescope (NGST) will provide valuable clues on the earliest episodes of galaxy formation from infra-red observations, although a relationship between the first galaxies and the properties of the IGM at high redshifts will not necessarily be easy to determine from such observations. Radio observations of the redshifted 21\\,cm line of neutral hydrogen might be capable of measuring the pre-reionization signal \\cite{mmr97,toz00,ili02,cgo02}, although such a measurement will be at the very edge of the capability of the next generation radio instruments such as Low Frequency Array (LOFAR).  Another possible channel to probe the pre-overlap stage of reionization is studying secondary CMB anisotropies \\cite{hw96,agh96,ksd98,gh98,jk98,pj98,hk99,h00,bru00,das00,ref00,swh01,sbp01,gj01}. While such a measurement is also in the future, observations of the secondary CMB anisotropies can provide constraints on the physical conditions in the IGM that are not accessible by other means. In fact, because CMB anisotropies are sensitive to the total Thompson optical depth, they can probe low levels of ionization in the mostly neutral gas - which cannot be done by radio observations of the redshifted 21\\,cm line.  Several scenarios have been proposed recently in which the universe undergoes a protracted episode of incomplete ionization prior to full reionization, either due to early supernova-driven winds \\cite{mfr01,oea01} or ionizations by energetic X-rays \\cite{oh01,vgs01}. This kind of ``smoldering'' reionization (which results in a partial reduction in the neutral hydrogen fraction) will be virtually impossible to detect with radio observations. Such a signature also cannot be detected by measuring just the spectrum of the fluctuations, because, as was shown in Gnedin \\& Jaffe \\shortcite{gj01}, the power spectrum of secondary anisotropies is only weakly dependent on the redshift of reionization. However, Gnedin \\& Jaffe \\shortcite{gj01} noticed that the fluctuations themselves were highly non-gaussian. In this paper we show how tests of the non-gaussianity of the secondary CMB anisotropies can detect the signature of an epoch of ``smoldering reionization''.  The physical reason behind such a possibility is simple: the earlier the epoch of ``smoldering reionization'' begins, the more nonlinear the objects that are responsible for the production of ionizing photons should be, and, therefore, the more non-gaussian the secondary CMB anisotropies will be. And it does not matter for this test what the sources of ionizations really are, only that they form early and, therefore, highly nonlinear at early times, when the linear fluctuations are smaller.  A big advantage of a morphological approach is that it is virtually independent of assumptions about the underlying cosmological model. For example, the redshifted 21cm emission and absorption will depend on the rate at which ionization front expand into the low density IGM, gobbling up the neutral gas on the way. And because secondary CMB anisotropies are dominated by nonlinear structures on small scales \\cite{gj01}, they are insensitive to the specific details of how ionization fronts expand. ",
        "conclusions": "We showed how the morphological analysis of the maps of the secondary CMB anisotropies on sub-arcminute angular scales can detect an extended period of ``smoldering'' reionization, during which the universe remains partially ionized. If the neutral hydrogen fraction during such a period is below about 50\\% but still well above $10^{-5}$, such a period will be detectable neither by radio observations of the redshifted 21cm line nor by IR observations of the Lyman-alpha forest. In that case morphology of the CMB anisotropies offers the best chance to probe the IGM at early times.  We computed each of six parameters as a function of the fractional area of the excursion set, $A_g$: 1) the number of regions in the excursion set, $N_{cl}$, 2) total perimeter in the excursion set, $P_g$, 3) genus of the excursion set defined as the number of regions minus the number of holes, $G_g$, 4) area of the largest (i.e. percolating) region, $A_p$, 5) perimeter of the largest region, $P_p$, and 6) the genus of the largest region, $G_p$. Three parameters ( $A_g$, $P_g$ and $G_g$) are known as the global Minkowski functionals of the excursion set, and $A_p$, $P_p$ and $G_p$ are the Minkowski functionals of largest (by area) region. We found $N_{cl}$, $P_g$, and $G_p$ are particularly sensitive to the non-gaussianity in $\\dt/T$ maps due to secondary reionization, they are also the most robust to effects of noise. $A_p$ and $P_p$ are not sensitive to this type of non-gaussianity, $G_g$ is significantly less sensitive then $N_{cl}$, $P_g$, and $G_p$. Using $N_{cl}$, $P_g$, and $G_p$ one can detect the non-gaussianity in the CMB maps with $S/N=1/2$ at the significance level of better than 99\\%.  Morphological analysis of the shapes of individual regions in the excursion set provides considerably more information about non-gaussianity of the maps and potentially may improve the current result both in terms of the resolution of the maps and signal to noise ratio. We reserve this study for the future work.  \\noindent {\\bf Acknowledgments:} S.Sh. acknowledges the support of the GRF 2002 grant at the University of Kansas. This work was partially supported by National Computational Science Alliance under grant AST-960015N and utilized the SGI/CRAY Origin 2000 array at the National Center for Supercomputing Applications (NCSA)."
    },
    "0206/astro-ph0206021_arXiv.txt": {
        "abstract": "The Time-of-Flight (TOF) system of the AMS detector gives the fast trigger to the read out electronics and measures velocity, direction and charge of the crossing particles.  The first version of the detector (called AMS-01) has flown in 1998 aboard of the shuttle Discovery for a 10 days test mission, and collected about $10^8$ events.  The new version (called AMS-02) will be installed on the International Space Station and will operate for at least three years, collecting roughly $10^{10}$ Cosmic Ray (CR) particles.  The TOF system of AMS-01 successfully operated during the test mission, obtaining a time resolution of 120 ps for protons and better for other CR ions.  The TOF system of AMS-02 will be different due to the strong fringing magnetic field and weight constraints.  [\\emph{Talk given at the ``First International Conference on Particle and Fundamental Physics in Space'', La Biodola, Isola d'Elba (Italy), 14 -- 19 May 2002.  To be published by Nuclear Physics B - Proceedings Supplement.}] \\vspace{1pc} ",
        "introduction": "The \\emph{Alpha Magnetic Spectrometer} (AMS)~\\cite{amsfirst} is a particle detector that will be installed on the International Space Station in 2005 to measure cosmic ray fluxes for at least three years.  Amongst the AMS goals we may cite: \\begin{itemize} \\item Search for cosmic antimatter;  \\item Search for dark matter signatures;  \\item Measurement of primary Cosmic Ray (CR) spectra below 1 TeV:  \\begin{itemize} \\item Hydrogen and Helium isotopes: solar modulation on a weekly basis;  \\item Very high statistics for CR ions below Iron;  \\item Precise measurements of electron and positron spectra;  \\item Cosmic gamma-rays spectrum. \\end{itemize} \\end{itemize}   During the precursor flight aboard of the shuttle Discovery (NASA STS-91 mission, 2--12 June 1998), AMS collected data for about 180 hours~\\cite{amsall}.  Figure~\\ref{AMS1} shows the test detector (called AMS-01 in the following), consisting of a permanent Nd-Fe-B magnet, six silicon tracker planes, a scintillator counters anticoincidence system, the time of flight (TOF) system consisting in four layers of scintillator counters and a threshold aerogel \\v{C}erenkov detector.  \\begin{figure*}[t] \\centering \\includegraphics[width=0.7\\textwidth]{\\ImgPath{ams1}} \\caption{The AMS detector for the STS-91 mission (AMS-01).}\\label{AMS1} \\end{figure*}  The TOF system of AMS-01~\\cite{tof1} was completely designed and built at the INFN Laboratories in Bologna, Italy. Its main goals are to provide the fast trigger to AMS readout electronics, and to measure the particle velocity ($\\beta$), direction, crossing position and charge.  In addition, it had to operate in space with severe limits for weight and power consuption (see section~\\ref{sec-tof2} for more details).  \\begin{figure*}[t] \\centering \\includegraphics[width=0.7\\textwidth]{\\ImgPath{ams2}} \\caption{The AMS detector to be installed on ISS (AMS-02).}\\label{AMS2} \\end{figure*}  Figure \\ref{AMS2} shows the new version of the detector (named AMS-02), that will be installed aboard of the ISS for a 3 years mission.  AMS-02 will be based on a superconducting magnet that will produce a dipolar field of 0.85 T maximum intensity (hence the separating power will be $BL^2 \\approx 0.85$ T \\um{m^2}).  In addition to a veto system of 8 organic scintillator paddles, a Silicon tracker of 8 $(x,y)$ planes will be placed inside the magnet. The new tracking system will reach a spatial resolution of $\\sigma_y = 10\\, \\um{{\\mu}m}$ on the bending plane and $\\sigma_x = 30\\, \\um{{\\mu}m}$ for the non-bending plane.  The total active area will be $\\approx 7 \\um{m^2}$ ($\\sim 2\\times10^5$ channels), and its rigidity resolution will be $\\Delta R/R \\approx 2\\%$ for $R = (1 \\div 30) \\um{GV}$ for protons.  The new time of flight system will consist of 4 planes of 8, 8, 10, 8 scintillator counters respectively (see section~\\ref{sec-tof2} below). Its time resolution will be $\\sim 140$ ps for protons and better for higher charged particles, making AMS-02 able to distinguish between negative and positive charged particles (hence between CR matter and antimatter) at the $10^{-9}$ level.  The main differences with respect to the TOF system of AMS-01 are due to the very strong magnetic field in the photomultiplier tubes (PMT) zone and to more severe weight limits.  Below the magnet, a proximity focusing RICH will substitute the AMS-01 threshold \\v{C}erenkov counter.  With a 2 cm thick aerogel radiator and a pixel plane 48 cm far from it, by reconstructing the \\v{C}erenkov angle it will reach a velocity resolution $\\Delta\\beta/\\beta \\sim 0.1\\%$ and by counting the emitted photons it will measure the charge of the incident particle.  This instrument will improve the AMS sensitivity to light elements isotopes up to $(12 \\div 13)$ GeV/A and will enhance element discrimination up to Fe.  Finally, two new instruments will be added to the AMS-02 detector with respect to AMS-01: a Transition Radiation Detector (TRD) on the top, and an electromagnetic calorimeter (ECAL) at the bottom of the detector.  By improving the discriminating power amongst electrons and protons up to high energies, TRD and ECAL will allow AMS-02 to measure with high statistics the positron and electron spectra up to $\\sim 300$ GeV, covering the most promising energy range for the search of supersymmetric dark matter particle annihilations. ",
        "conclusions": "The time of flight system of AMS-02 will have the same goals of the TOF system of AMS-01, but it will operate with more severe conditions.  The low weight and power consumption budgets and the very strong fringing magnetic field in the phototubes zone will cause a worsening of the time resolution, with respect to AMS-01.  On the other hand, the new read-out electronics will make the new system able to reach a better charge resolution than the previous one, and the trigger efficiency will be unaffected by those problems.  In addition, fast additional information based on the energy release of the crossing particles will be sent to the trigger logics, in order to be able to flag CR nuclei with charge greater than two."
    },
    "0206/astro-ph0206398_arXiv.txt": {
        "abstract": "An imaging survey of CO(1$-$0), HCN(1$-$0), and HCO$^+$(1$-$0) lines in the centers of nearby Seyfert galaxies has been conducted using the Nobeyama Millimeter Array and the RAINBOW interferometer. Preliminary results reveal that 3 Seyferts out of 7 show abnormally high HCN/CO and HCN/HCO$^+$ ratios, which cannot occur even in nuclear starburst galaxies. We suggest that the enhanced HCN emission originated from X-ray irradiated dense obscuring tori, and that these molecular line ratios can be a new diagnostic tool to search for ``pure'' AGNs. According to our HCN diagram, we suggest that NGC 1068, NGC 1097, and NGC 5194 host ``pure'' AGNs, whereas Seyfert nuclei of NGC 3079, NGC 6764, and NGC 7469 may be ``composite'' in nature. ",
        "introduction": "Dense molecular matter is considered to play various roles in the vicinity of active galactic nuclei (AGNs). The presence of dense and dusty interstellar matter (ISM), which obscures the broad line regions in AGNs, is inevitable at a few pc - a few 10 pc scale according to the unified model of Seyfert galaxies. This circumnuclear dense ISM could be a reservoir of fuel for nuclear activity, and also be a site of massive star formation.  In order to investigate dense molecular matter in the centers of Seyfert galaxies, we have conducted an imaging survey of CO(1$-$0), HCN(1$-$0), and HCO$^+$(1$-$0) lines in nearby Seyfert galaxies using the Nobeyama Millimeter Array (NMA). High resolution HCN observations of Seyfert galaxies are of interest because unusually strong HCN emission has been reported in the type-2 Seyfert galaxies NGC 1068 (Jackson et al. 1993; Tacconi et al. 1994; Helfer \\& Blitz 1995) and NGC 5194 (Kohno et al.\\ 1996). The HCN/CO integrated intensity ratios in brightness temperature scale, $R_{\\rm HCN/CO}$ hereafter, within the central $r \\sim$ a few 10 pc region exceed 0.4, which {\\it is never observed in non-Seyfert galaxies} including nuclear starburst galaxies ($R_{\\rm HCN/CO} < 0.3$; see Kohno et al.\\ 1999 and references therein). ",
        "conclusions": "\\begin{figure} \\plotone{Kohno-fig4-n1097b.eps} \\caption{ CO and HCN intensity profiles through the nucleus along P.A. = 60$^\\circ$ line (left), azimuthally averaged radial distributions of CO and HCN (top right), and HCN/CO ratio (bottom right). } \\end{figure}  We have observed extremely strong HCN emission in 3 Seyferts out of 7. What is the nature of these ``HCN-enhanced Seyferts''? Here we compare the observed line ratios in Seyferts with those in nuclear starburst galaxies, which were also measured with similar angular resolutions. In Figure 5, it is immediately evident that Seyferts without abnormal HCN enhancements, i.e. NGC 3079, NGC 6764, and NGC 7469, show $R_{\\rm HCN/CO}$ and $R_{\\rm HCN/HCO^+}$ values just comparable to those in nuclear starbursts; they have $R_{\\rm HCN/CO}$ less than 0.3, and $R_{\\rm HCN/HCO^+}$ ranging from 0.5 to 1.5. On the other hand, HCN-enhanced Seyferts, i.e. NGC 1068 and NGC 1097, also have very high $R_{\\rm HCN/HCO^+}$ values ($> 2$). Note that Nguyen-Q-Rieu et al. (1992) reported a very high $R_{\\rm HCN/HCO^+}$ in NGC 3079 and Maffei 2 ($>3$), yet our new simultaneous measurements gave moderate ($\\sim 1$) ratios.  We propose that these two groups in our ``HCN diagram'' (Figure 5) can be understood in terms of ``AGN - nuclear starburst connection'' (note that this should not be confused with ``AGN - starburst cohabitation'', which often refers to the association of AGN with star formation on galactic scales in AGN hosts). In the Seyferts with line ratios comparable to those in nuclear starburst galaxies, it seems likely that nuclear starburst (presumably in the dense molecular torus) is associated with the Seyfert nucleus (i.e., ``composite''). In the nuclear regions of composite Seyferts, HCO$^+$ fractional abundance is expected to increase due to frequent supernova (SN) explosions. In fact, in evolved starbursts such as M82, where large scale outflows have occurred due to numerous SN explosions, HCO$^+$ is often stronger than HCN (e.g. Nguyen-Q-Rieu et al.\\ 1992). On the other hand, the HCN-enhanced Seyferts, which shows $R_{\\rm HCN/CO} > 0.3$ and $R_{\\rm HCN/HCO^+} > 2$, would host ``pure'' AGNs, where there is no associated nuclear starburst activity. In such a condition, the HCN line can be very strong because it has been predicted that fractional abundance of HCN is enhanced by strong X-ray radiation from AGN (Leep \\& Dalgarno 1996), resulting in abnormally high $R_{\\rm HCN/CO}$ and $R_{\\rm HCN/HCO^+}$ values. Our interpretation is supported by other wavelength data; for instance, NGC 1068 has been claimed as a pure Seyfert (Cid Fernandes et al.\\ 2001 and references therein), whereas NGC 6764 (Schinnerer et al.\\ 2000) and NGC 7469 (Genzel et al.\\ 1995) have a composite nature. We need further analysis to validate the proposed interpretation, but if it is the case, this will serve as a new way to investigate the nature of AGNs; although this technique requires high angular resolution observations in order to avoid contaminations from extended circumnuclear star-forming regions, it has some advantages (e.g., not being affected by dust extinction).     \\begin{figure} \\plotone{Kohno-fig5-diagram.eps} \\caption{ Molecular line ratios in Seyfert and non-Seyfert (nuclear starburst) galaxies. } \\end{figure}"
    },
    "0206/astro-ph0206459_arXiv.txt": {
        "abstract": "{ We present the results of three-dimensional hydrodynamical simulations of the subsonic thermonuclear burning phase in type Ia supernovae. The burning front model contains no adjustable parameters so that variations of the explosion outcome can be linked directly to changes in the initial conditions. In particular, we investigate the influence of the initial flame geometry on the explosion energy and find that it appears to be weaker than in 2D. Most importantly, our models predict global properties such as the produced nickel masses and ejecta velocities within their observed ranges without any fine tuning. ",
        "introduction": "In a series of papers \\citep{reinecke-etal-99a, reinecke-etal-99b, reinecke-etal-02a} we described a new numerical tool to simulate thermonuclear explosions of white dwarfs in two and three spatial dimensions. Our aim was to construct models of type Ia supernova (SN Ia) explosions that are as free from non-physical parameters as currently feasible. Thermonuclear burning, in particular, is represented by a subsonic turbulent flame whose local velocity is derived from a subgrid-scale model for unresolved turbulent fluctuations. Solved in combination with the compressible Euler equations, this model contains no free parameters that could be adjusted in order to fit SN Ia observations. Consequently, the initial white dwarf model (composition, central density, and velocity structure), as well as assumptions about the location, size and shape of the flame surface as it first forms fully determine the simulation results. Here, we concentrate on variations of the latter.  In \\cite{reinecke-etal-99a}, we confirmed the earlier result of \\cite{niemeyer-etal-96} that, at least in 2D, the explosion energy and amount of $^{56}$Ni produced in the explosion (which determines the brightness of an SN Ia) are sensitive to the ignition conditions. To be more precise, a more complicated topology of the initial nuclear flame seems to lead to higher Ni-production and, consequently, more powerful explosions. One might even speculate that the randomness of the ignition process could be the reason for the observed spread in properties of normal SNe Ia \\citep{hillebrandt-niemeyer-2000}.  This article continues the presentation of numerical simulations of SNe Ia by testing the effect of different initial conditions on the simulation outcome in three dimensions. In this context, the simultaneous runaway at several different spots in the central region of the progenitor star is of particular interest. A plausible ignition scenario was suggested by \\cite{garcia-woosley-95}.  All simulations were performed using the algorithms presented by \\cite{reinecke-etal-02a, reinecke-etal-99a}. The initial model for the white dwarf, assumed to be at the Chandrasekhar mass, as well as the grid geometry and symmetry assumptions, are identical to the setup described in Sects.\\ 3 and 4 of \\cite{reinecke-etal-02a}.  Sect.\\ \\ref{multipoint} of this paper presents the initial front geometry and its temporal evolution for two multi-point ignition calculations. Various aspects of all three-dimensional calculations are then compared in Sect.\\ \\ref{seccomp3d}. As far as possible, this comparison is then extended to other existing results of SN Ia simulations. In the case of parameterized one-dimensional calculations such comparisons are difficult, since most of the published data have no analogy in three dimensions. As an alternative, we suggest to track and compare the amount of burned material as function of the density at which the burning took place.  Generally speaking, we confirm the earlier 2D results: models with more, or better resolved, ignition spots tend to produce more radioactive Ni also in 3D, although the effect is somewhat smaller. But together with the gain of Ni by going from 2D to 3D our models predict amounts that are in good agreement with those inferred from the observations \\citep{contardo-etal-00}. This will also be discussed in Sect.\\ \\ref{seccomp3d}, and our conclusions follow in Sect.\\ \\ref{discuss}.  \\section {Multi-point ignition scenarios} \\label{multipoint}  \\begin{figure}[tbp] \\centerline{\\includegraphics[width=0.48\\textwidth]{h3657_1.eps}} \\caption{Snapshots of the flame front for scenario b5\\_3d. The fast merging between the leading and trailing bubbles and the rising of the entire burning region is clearly visible. One ring on the coordinate axes corresponds to $10^7$cm.} \\label{b5_3d_front} \\end{figure} \\begin{figure}[tbp] \\centerline{\\includegraphics[width=0.48\\textwidth]{h3657_2.eps}} \\caption{Snapshots of the flame front for the highly resolved scenario b9\\_3d. One ring on the coordinate axes corresponds to $10^7$cm.} \\label{b9_3d_front} \\end{figure}  Two different calculations were performed to investigate the off-center ignition model in three dimensions. The simulation b5\\_3d\\_256 was carried out on a grid of 256$^3$ cells with a central resolution of 10$^6$\\,cm and contained five bubbles with a radius of $3\\cdot 10^6$\\,cm, which were distributed randomly in the simulated octant within $1.6\\cdot10^7$\\,cm of the star's center. As an additional constraint, the bubbles were required not to overlap significantly with the other bubbles or their mirror images in the other octants. This maximizes the initial flame surface. The exact algorithm for the positioning of the bubbles was the following: First, the centers of the bubbles were chosen randomly within the radius of $1.6\\cdot10^7$\\,cm. A realization was accepted when the distance between any two bubble centers (including the mirror images across the coordinate planes) was larger than 1.3 bubble radii. A particular realization of these initial conditions and its temporal evolution is shown in Fig.\\ \\ref{b5_3d_front}. The locations of the bubble centers are listed in Table \\ref{b5_table}.  \\begin{table} \\begin{center} \\begin{tabular}{|c|r|r|r|r|r|} \\hline Part.\\ No.& 1&2&3&4&5 \\\\ \\hline x-Pos. & 128 & 35 & 42 & 85 & 63 \\\\ \\hline y-Pos. & 76 & 60 & 129 & 27 & 85 \\\\ \\hline z-Pos. & 27 & 30 & 30 & 41 & 83 \\\\ \\hline \\end{tabular} \\end{center} \\caption{Approximate position of the bubble centers in the model b5\\_3d\\_256. All lengths are given in kilometers.} \\label{b5_table} \\end{table}  In a way very similar to the evolution of earlier two-dimensional simulations (cf.\\ Fig.\\ 6 of \\citealt{reinecke-etal-99b}) the flame kernels closer to the center are elongated very rapidly and connect to the outermost bubbles within 0.15\\,s. The whole burning region disconnects from the coordinate planes and starts to float slowly towards the stellar surface.  In an attempt to reduce the initially burned mass as much as possible without sacrificing too much flame surface, the very highly resolved model b9\\_3d\\_512 was constructed. It contains nine randomly distributed, non-overlapping bubbles with a radius of 2$\\cdot$10$^6$\\,cm within $1.6\\cdot10^7$\\,cm of the white dwarf's center. To properly represent these very small bubbles, the cell size was reduced to $\\Delta=5\\cdot 10^5$\\,cm, so that a total grid size of $512^3$ cells was required.  Starting out with very small flame kernels is desirable, because the initial hydrostatic equilibrium is better preserved if only a little mass is burned instantaneously. Furthermore the floating bubbles are expected to be even smaller in reality than in the presented model ($r\\lessapprox5\\cdot10^5$\\,cm, see \\citealt{garcia-woosley-95}).  \\begin{table} \\begin{center} \\begin{tabular}{|c|r|r|r|r|r|} \\hline Part.\\ No.& 1&2&3&4&5 \\\\ \\hline x-Pos. & 70.2 & 106.4 & 22.8 & 77.8 & 137.2 \\\\ \\hline y-Pos. & 48.6 & 24.6 & 121.2 & 106.7 & 68.2 \\\\ \\hline z-Pos. & 128.0 & 105.2 & 77.4 & 88.3 & 20.4 \\\\ \\hline\\hline Part.\\ No.& 6&7&8&9& \\\\ \\hline x-Pos. & 47.7 & 41.4 & 27.5 & 42.2 & \\\\ \\hline y-Pos. & 45.5 & 133.1 & 24.1 & 85.0 & \\\\ \\hline z-Pos. & 91.2 & 34.0 & 39.3 & 21.7 & \\\\ \\hline \\end{tabular} \\end{center} \\caption{Position of the bubble centers in the model b9\\_3d\\_512. All lengths are given in kilometers.} \\label{b9_table} \\end{table}  The center locations of our particular realization are given in Table \\ref{b9_table}. Snapshots of the front evolution (Fig.\\ \\ref{b9_3d_front}) exhibit features very similar to those observed in the lower-resolution 3D models. Only in the last plot the formation of additional small-scale structures becomes evident.  \\begin{figure*}[tbp] \\centerline{\\includegraphics[width=0.8\\textwidth]{h3657_3.eps}} \\caption{Time evolution of the chemical composition in the models b5\\_3d\\_256 and b9\\_3d\\_512.} \\label{nuc_evolve3d} \\end{figure*}  The direct comparison of the nuclear evolution with model b5\\_3d\\_256 in Fig.\\ \\ref{nuc_evolve3d} nevertheless reveals differences in the early explosion stages: the high-resolution simulation produces slightly more nickel and considerably more $\\alpha$-particles during the first half second. This phenomenon is most likely explained by the discrepancy in the ratio between the initial flame surface and the volume of burned material in the two models. The higher burned mass in the five-bubble model initiates an early bulk expansion of the star and therefore causes a rapid drop of the central density. Since its flame surface is quite small compared to the burned volume, only relatively little mass can be burned at high densities. In the nine-bubble model, on the other hand, the star expands more slowly at first because less material is burned instantaneously, and the mass fraction of $\\alpha$-particles in the reaction products is consequently rather high. Since the energy buffered in those $\\alpha$-particles is not immediately used to drive the expansion, the flame can consume more material at higher densities and has more time to increase its surface as a result of hydrodynamical instabilities. ",
        "conclusions": "\\label{discuss}  In this paper we have presented the results of three-dimensional numerical simulations of thermonuclear deflagration fronts in Chandrasekhar mass white dwarfs composed of equal amounts of carbon and oxygen. We could show that independent of the details of the ignition process (which is still far from being well understood) the white dwarf is always disrupted by the release of nuclear energy. As far as we could check at present, the properties of our models are in good agreement with observations of typical type Ia supernovae. In particular the explosion energy and the average chemical composition of the ejecta seem to fit the observations (see also Table 1). This success of the models was obtained without introducing any non-physical parameters, but just on the basis of a physical and numerical model of subsonic turbulent burning fronts. We also stress that our models give clear evidence that the often postulated deflagration-detonation transition is not needed to produce sufficiently powerful explosions.   \\begin{table}[htbp] \\centerline{ \\begin{tabular}{|l|c|c|c|} \\hline model name & $m_{\\text{Mg}}$ [$M_\\odot$]&$m_{\\text{Ni}}$ [$M_\\odot$]&$E_{\\text{nuc}}$ [$10^{50}$\\,erg\\vphantom{\\raisebox{2pt}{\\large A}}] \\\\ \\hline c3\\_3d\\_256 & 0.177 & 0.526 & 9.76 \\\\ \\hline b5\\_3d\\_256 & 0.180 & 0.506 & 9.47 \\\\ \\hline b9\\_3d\\_512 & 0.190 & 0.616 & 11.26\\phantom{0} \\\\ \\hline \\end{tabular} } \\caption{Overview over element production and energy release of all discussed supernova simulations} \\label{burntable} \\end{table}  There are certainly several desirable additions and improvements. The most crucial question still seems to be the ignition process, although in 3D the dependence of the final outcome is weaker than in our previous 2D models. In principle one could try to simulate the ignition phase with the numerical models we have developed, but because of the much longer time scales this requires huge amounts of computer time. First attempts in this direction are presently under way (cf. \\citealt{hoeflich-stein-2002}).  An improvement of the combustion model already mentioned in \\cite{reinecke-etal-99a}, i.e. the full reconstruction of all thermodynamic quantities from their mean values in every grid cell cut by the burning front, has now been completed and was applied to the Landau-Darrieus instability \\citep{roepke-etal-02}. This new model should be implemented into the full code which, in principle, seems to be possible but difficult. In any case, this would increase the predictive power and reliability of the models.  In a similar direction, it might be worthwhile to test other subgrid models in combination with our level set method. We do not expect major changes because if another subgrid-model would give us higher (or lower) burning speed on the grid scale this would lead to more (or less) damping of small-scale structures leaving the product of the two (and therefore the fuel consumption rate) approximately unchanged. However, again, such studies are presently in progress.  To conclude, we think that we have made another step towards the understanding of type Ia supernovae."
    },
    "0206/astro-ph0206490_arXiv.txt": {
        "abstract": "{ We find evidence that the two high frequency QPOs in Sco X-1 are, more often than not,  approximately in the 2:3  frequency ratio familiar from studies of black hole candidates (e.g., XTE J1550-564,  Remillard et al. 2002). This implies that the double kHz QPO phenomenon in neutron stars has its origin in properties of strong-field gravity and has little to do with the rotation of a stellar surface or any magnetic field structure anchored in the star. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206173_arXiv.txt": {
        "abstract": "We consider the effect of inhomogeneous neutrino degeneracy on Big Bang nucleosynthesis for the case where the distribution of neutrino chemical potentials is given by a Gaussian.  The chemical potential fluctuations are taken to be isocurvature, so that only inhomogeneities in the electron chemical potential are relevant. Then the final element abundances are a function only of the baryon-photon ratio $\\eta$, the effective number of additional neutrinos $\\Delta N_\\nu$, the mean electron neutrino degeneracy parameter $\\bar \\xi$, and the rms fluctuation of the degeneracy parameter, $\\sigma_\\xi$. We find that for fixed $\\eta$, $\\Delta N_\\nu$, and $\\bar \\xi$, the abundances of $^4$He, D, and $^7$Li are, in general, increasing functions of $\\sigma_\\xi$.  Hence, the effect of adding a Gaussian distribution for the electron neutrino degeneracy parameter is to {\\it decrease} the allowed range for $\\eta$.  We show that this result can be generalized to a wide variety of distributions for $\\xi$. ",
        "introduction": "Many modifications to the standard model of Big Bang nucleosynthesis (BBN) have been explored \\cite{malaney}. One of the most exhaustively investigated variations on the standard model is neutrino degeneracy, in which each type of neutrino is allowed to have a non-zero chemical potential \\cite{wagoner}, and a number of models have been proposed to produce a large lepton degeneracy \\cite{gelmini}-\\cite{murayama}. More recently, observations of the cosmic microwave background (CMB) fluctuations have been combined with BBN to further constrain the neutrino chemical potentials \\cite{lesg}-\\cite{orito2}.  An interesting variation on these models is the possibility that the neutrino degeneracy is inhomogeneous \\cite{dolgov1}-\\cite{dibari}.  The consequences of inhomogeous neutrino degeneracy for BBN were examined by Dolgov and Pagel \\cite{dolgov2} and Whitmire and Scherrer \\cite{whitmire}. Dolgov and Pagel examined models in which the length scale of the inhomogeneity was sufficiently large to produce an inhomogeneity in the presently-observed abundances of the elements produced in BBN. Whitmire and Scherrer investigated inhomogeneities in the neutrino degeneracy on smaller scales; in these models the element abundances mix to produce a homogeneous final element distribution. Using a linear programming technique, they derived upper and lower bounds on the baryon-to-photon ratio $\\eta$ for {\\it arbitrary} distributions of the neutrino chemical potentials and showed that the upper bound on $\\eta$ could be considerably relaxed. However, the resulting distributions for the neutrino chemical potentials were quite unnatural.  Hence, in this paper, we examine a more restricted class of models, in which the distribution of the chemical potentials is taken to be a Gaussian.  In the next section, we discuss our model for inhomogeneous neutrino degeneracy. We calculate the effect of these inhomogeneities on the final element abundances and discuss our results in Sec. 3. We find that, in most cases, the effect of Gaussian inhomogeneities in the electron neutrino chemical potential is to increase the abundances of deuterium, $^4$He, and $^7$Li relative to their abundances in models with homogeneous neutrino degeneracy. ",
        "conclusions": "The model described in the previous section can be completely specified by two parameters, the (inhomogeneous) electron neutrino degeneracy parameter, $\\xi_e$, and the additional (homogeneous) energy density due to the degeneracy of all three neutrinos plus any additional relativistic component.  We parametrize the latter in terms of $\\Delta N_\\nu$, the effective number of additional neutrinos.  This second parameter hides our ignorance about the compensation mechanism and about the degeneracies among the other two types of neutrinos. In our simulation, we take $\\xi_e$ to be homogeneous within a given horizon volume during nucleosynthesis. Different horizon volumes may have different values of $\\xi_e$, which are given by the distribution function $f(\\xi_e)$, i.e, the probability that a given horizon volume has a value of $\\xi_e$ between $\\xi_e$ and $\\xi_e + d\\xi_e$.  (Since we are considering only inhomogeneities in electron neutrinos, we now drop the $e$ subscript). We take $f(\\xi)$ to have a Gaussian distribution with mean $\\bar \\xi$ and rms fluctuation $\\sigma_\\xi$: \\begin{equation} f(\\xi) = {1 \\over \\sqrt {2 \\pi}\\sigma_\\xi} \\exp[- (\\xi - \\bar \\xi)^2/ 2\\sigma_\\xi^2] \\end{equation} Then the final primordial element abundances, for a fixed value of $\\eta$ and $\\Delta N_\\nu$, will be functions of $\\bar \\xi$ and $\\sigma_\\xi$; we can write, for a given nuclide $A$, \\begin{equation} \\label{XA} \\bar X_A = \\int_{-\\infty}^\\infty X_A(\\xi)f(\\xi) d\\xi, \\end{equation} where $X_A(\\xi)$ is the mass fraction of $A$ as a function of $\\xi$, and $\\bar X_A$ is the mass fraction of $A$ averaged over all space; after the matter is thoroughly mixed, $\\bar X_A$ will be the final primordial mass fraction.  A full treatment for all possible values of $\\eta$, $\\Delta N_\\nu$, $\\bar \\xi$, and $\\sigma_\\xi$ is impractical.  We have chosen to concentrate on variations in the latter two quantities, since we are most interested in the effects of inhomogeneities in the chemical potential. Because of the large number of free parameters and the difficulty of exhaustively searching all of parameter space, our goal is to discern any general results which are independent of $\\eta$ and $\\Delta N_\\nu$.  There are now strong limits on $\\eta$ from the cosmic microwave background alone, independent of BBN.  We examine two extreme values for $\\eta$: $\\eta = 4 \\times 10^{-10}$ and $\\eta = 1 \\times 10^{-9}$; these represent very conservative lower and upper bounds on $\\eta$ from the CMB in models with non-zero neutrino degeneracy \\cite{Kneller}.  For $\\Delta N_\\nu$, we consider $\\Delta N_\\nu = 0$ and 5.  Note that the first of these is only possible if the extra energy density in the degenerate electron neutrinos is compensated by a decrease in the energy density in some other relativistic component. For each of these cases, we calculate the abundances of $^4$He, D, and $^7$Li as a function of $\\sigma_\\xi$ for $\\bar \\xi = -1$ to $+1$ in steps of 0.5.  Our results are displayed in Figures 1-3.  In each of these figures, we also show observational limits on the primordial element abundances from Ref. \\cite{olive}: $2.9 \\times 10^{-5} < ({\\rm D/H}) < 4.0 \\times 10^{-5}$, $0.228 < Y_P < 0.248$, and $-9.9 < \\log(^7{\\rm Li/H}) < -9.7$.  The general behavior of the element abundances in Figs. 1-3 is very clear.  As expected, for $\\sigma_\\xi \\ll \\bar \\xi$, the abundances of deuterium, $^4$He, and $^7$Li are unchanged from their values in the corresponding homogeneous model with the same value of $\\bar \\xi$.  At the opposite limit, when $\\sigma_\\xi \\gg \\bar \\xi$, the models all converge to a single limiting value; again, this is what one would naively expect.  What is interesting is that, with a few exceptions, the introduction of a Gaussian distribution of values for $\\xi$ results in an {\\it increase} in the abundance of each element relative to the corresponding homogeneous model with the same value of $\\bar \\xi$.  The only exceptions occur for $^4$He with negative values of $\\bar \\xi$, (for which $Y_P$ is far too large to be physically reasonable), and some of the $^7$Li curves, for which there is a tiny decrease in the $^7$Li abundance over a short range of $\\sigma_\\xi$ values.  This result may seem surprising, but it is a simple consequence of the behavior of $X_A(\\xi)$.  In particular, if $X_A(\\xi)$ is a convex function ($X_A^{\\prime \\prime}(\\xi) > 0$), then Jensen's inequality \\cite{feller} gives \\begin{equation} \\label{jensen} \\int_{-\\infty}^\\infty X_A(\\xi)f(\\xi) d\\xi > X_A(\\bar \\xi). \\end{equation} We find, for example, for $\\Delta N_\\nu = 0$, and both values of $\\eta$, that our $X(\\xi)$ curves are all convex in the range $-2 < \\xi < 2$, with the exception of $^4$He at $\\xi < -1$, and $^7$Li with $\\eta = 1 \\times 10^{-9}$.  These are precisely the regimes for which we observe equation (\\ref{jensen}) to fail. Of course, none of the $X_A(\\xi)$ curves is convex for all values of $\\xi$; the practical condition for equation (\\ref{jensen}) to hold is that the $X(\\xi)$ curves be convex as long as $f(\\xi)$ is non-negligible.  \\onecolumn  \\begin{figure} \\begin{center} \\epsfysize=8.0truein \\epsfbox{fig1.eps} \\vspace*{0.5truecm} \\end{center} \\vspace*{-0.3truecm} \\caption{The primordial $^4$He mass fraction, $Y_P$, as a function of the rms fluctuation in the electron chemical potential $\\sigma_\\xi$, for the indicated value of the mean electron neutrino chemical potential $\\bar \\xi$.  Each figure corresponds to the indicated value of the baryon-photon ratio $\\eta$ and the effective number of extra neutrinos $\\Delta N_\\nu$.  The gray shaded region gives observational limits on $Y_P$ from Ref. [19].} \\label{fig1} \\end{figure}  \\begin{figure} \\begin{center} \\epsfysize=8.0truein \\epsfbox{fig2.eps} \\end{center} \\baselineskip 7pt \\caption{As Fig. 1, for the primordial ratio of deuterium to hydrogen, (D/H).} \\label{fig2} \\end{figure}  \\begin{figure} \\begin{center} \\epsfysize=8.0truein \\epsfbox{fig3.eps} \\end{center} \\caption{As Fig. 1, for the primordial ratio of $^7$Li to hydrogen, ($^7$Li/H).} \\label{fig3} \\end{figure}  \\twocolumn  This simple behavior allows us to draw some useful general conclusions.  In models in which $\\xi$ and $\\Delta N_\\nu$ are allowed to vary freely, if we fix $\\xi$ and trace out the allowed region in the $\\eta$, $\\Delta N_\\nu$ plane, then the upper and lower bounds on $\\eta$ are set primarily by the upper observational bound on $^7$Li and the upper observational limit on D, respectively, with the $^4$He limits serving primarily to set the bounds on $\\Delta N_\\nu$ \\cite{Kneller}. However, our results indicate that the general effect of going from a homogeneous to an inhomogeneous distribution in $\\xi$ is to {\\it increase} both the deuterium and the $^7$Li abundances.  (Again, we note a slight decrease in $^7$Li over a small range in $\\sigma_\\xi$, but this is a tiny effect).  Hence, the net effect of introducing this inhomogeneity will be to {\\it decrease} the allowed range for $\\eta$, in comparison with the corresponding homogeneous model.  This is a rare example in the study of BBN in which the introduction of an extra degree of freedom does nothing to increase the allowed range for $\\eta$. Instead, the effect of adding a Gaussian distribution of values for $\\xi$ is to decrease the allowed range for $\\eta$.  Although we have assumed a Gaussian distribution for $\\xi$, our results are much more general. In particular, as long as our distribution $f(\\xi)$ is negligible over the range of values of $\\xi$ for which $X_A(\\xi)$ is not a convex function, we expect equation (\\ref{jensen}) to hold. This would apply, for example, to a top hat distribution with the same values of $\\sigma_\\xi$ as those examined here. Moreover, the distribution $f(\\xi)$ need not even be symmetric for our results to apply.  Our results contrast with those of Ref. \\cite{whitmire}, which found an expanded upper limit on $\\eta$ in models with inhomogeneous $\\xi$. The reason for this difference is that the models examined in Ref. \\cite{whitmire} allowed for an arbitrary distribution in $\\xi$, and large increases in $\\eta$ occurred for bizarre distributions in $\\xi$.  In particular, the distributions in Ref. \\cite{whitmire} sampled extreme values for $\\xi$, outside the range for which all of the $X_A(\\xi)$ functions are convex."
    },
    "0206/astro-ph0206035_arXiv.txt": {
        "abstract": "We examine the theoretical implications of a population of low-mass helium-core white dwarfs in globular clusters. In particular, we focus on the observed population in the core of NGC~6397, where several low-mass white dwarf canditates have been identified as ``non-flickerers'' by Cool and collaborators. Age and mass estimates from cooling models, combined with dynamical and evolutionary considerations, lead us to infer that the dark binary companions are C/O white dwarfs rather than neutron stars. Furthermore, we find that the progenitor binaries very likely underwent an exchange interaction within the last $10^9$~years. We examine the prospects for detecting a similar population in other globular clusters, with particular attention to the case of 47~Tuc. ",
        "introduction": "Binaries play a central role in the internal dynamical processes that drive globular cluster evolution (e.g., Hut et al.\\ 1992; Bailyn 1995). Of particular importance is the primordial binary star population (Heggie 1975; Hut et al.\\ 1992), which provides a crucial internal energy source for the cluster by virtue of inelastic scattering encounters.  Some of the more dramatic products of these encounters include recycled pulsars (Rappaport, Putney, \\& Verbunt 1989; Phinney 1996), low-mass X-ray binaries (Verbunt \\& Johnston 1996), cataclysmic variables (Di Stefano \\& Rappaport 1994; Cool et al. 1995; Grindlay et al. 2001) and blue stragglers (Sigurdsson, Davies, \\& Bolte 1994; Lombardi, Rasio, \\& Shapiro 1996).  Cool et al.\\ (1998) and Edmonds et al.\\ (1999) (hereafter CGC and EGC, respectively) have reported the detection of a new stellar population near the center of the core-collapsed globular cluster NGC~6397. They named them ``non-flickerers'' and tentatively identified them as low-mass helium-core white dwarfs (hereafter HeWDs). Several similar, but fainter, objects have been discovered by Taylor et al. (2001) in the same cluster. HeWDs are the result of the truncated evolution of low- and intermediate-mass stars in binaries (Kippenhahn, Kohl \\& Weigert 1967), in which the hydrogen envelope of an evolving star is removed before the degenerate core is massive enough to burn helium to carbon. Such objects are a generic by-product of mass transfer in close binaries with evolved low-mass donor stars and may offer insight into the formation history of the compact object population and its coupling to the global cluster evolution.  In this paper we use white dwarf cooling models, together with simple binary evolution and dynamical models, to examine the nature of the ``non-flickerers,'' extending the initial discussion of EGC. We confirm that they must be HeWDs with more massive, dark binary companions. We place these systems in the broader dynamical context appropriate to the underlying binary population and examine the results as a function of globular cluster parameters, with particular application to NGC~6397 and 47~Tuc.  The observations and their immediate implications are briefly summarized in \\S2. In \\S 3 we discuss the formation and cooling of HeWDs. Applications to NGC~6397 and 47~Tuc are presented in \\S4 and \\S5, respectively, while other clusters are discussed briefly in \\S6. ",
        "conclusions": "\\label{conclusions}  We have examined the various possible evolutionary pathways for the formation of NF systems, identified as HeWD by CGC, in globular cluster cores. We have payed particular attention to the cluster NGC~6397 where these objects were first discovered. We find that most likely the NGC~6397 NFs are HeWD with CO core white dwarfs as their dark companions. This model satisfies all constraints imposed on the lifetimes and masses of white dwarfs (inferred from the observations and the dynamical properties of the host cluster). The low masses of CO white-dwarf companions (relative to neutron-star companions) allow recent common-envelope evolution that produces (i) orbits tight enough to avoid disruption due to dynamical interactions, and (ii) young, bright HeWD to explain their position in the CMD. Neutron-star dark companions would be too massive and would result either in stable mass transfer (producing systems that are too wide to survive exchange interactions for a significant time in a dense cluster core), or in HeWD too old and too faint.  The common envelope episode also results in a moderate mass hydrogen envelope on the white dwarf surface, leading to an estimate of the cooling age that is consistent with the dynamical considerations. A massive hydrogen envelope, in which nuclear burning is important, would imply that the observed systems are older than they are likely to be, based on their probability of survival in the cluster core.  We have extended our analysis of NGC~6397 to other globular clusters to illustrate how the production of NFs depends on cluster environment. In particular, we have studied the potential NF population in 47~Tuc, which is both less dense and more massive than NGC~6397, and for which a wealth of observational data is becoming available. In this cluster surviving NFs with neutron star companions are possible, but they will be older and more massive (and correspondingly fainter) than those with C/O white dwarf companions. This is again the consequence of whether binaries can survive intact, otherwise being transformed by exchange interactions or gravitational wave-induced mergers.  We anticipate that these models can be extended to other clusters as observational results accrue. The inferred white dwarf cooling ages will allow us to place constraints on the exchange interaction history of clusters in a manner that is independent of other methods, such as the properties of millisecond pulsar binaries."
    },
    "0206/astro-ph0206086_arXiv.txt": {
        "abstract": "Studies of disks around young brown dwarfs are of paramount importance to our understanding of the origin, diversity and early evolution of sub-stellar objects. Here we present first results from a systematic search for disk emission in a spectroscopically confirmed sample of young objects near or below the sub-stellar boundary in a variety of star-forming regions. Our VLT and Keck $L^{\\prime}$-band observations of the $\\sigma$ Orionis and TW Hydrae associations suggest that if a majority of brown dwarfs are born with disks, at least the inner regions of those disks evolve rapidly, possibly clearing out within a few million years. ",
        "introduction": "The current paradigm of low-mass star formation holds that a young star accretes material from a circumstellar disk during the first (few) million years of its lifetime. The disk also provides the building material for planets. Over the past two decades, substantial observational evidence has accumulated to support this picture (Mannings, Boss \\& Russell 2000). However, much of that evidence rests on studies of stars within a relatively narrow mass range. In particular, there are few observational constraints on the formation of objects near or below the sub-stellar boundary.  If a large fraction of young brown dwarfs indeed harbor disks, the implication is that extremely low-mass objects may form via a mechanism similar to higher mass stars. Thus, reliable determination of the disk frequency as a function of age in young sub-stellar populations is critical to our understanding of their origin, diversity and early evolution.  We have commenced a program to obtain $L^{\\prime}$-band data on a large, {\\it spectroscopically confirmed} sample of young objects near or below the sub-stellar boundary in a number of nearby star-forming regions. $L^{\\prime}$-band photometry is much better at detecting disk excess above the photospheric emission, and is less susceptible to the effects of disk geometry, than measurements at shorter wavelengths. ",
        "conclusions": "Of the six $\\sigma$ Ori sources, only one (SOri 12) shows significant $K$-$L^{\\prime}$ excess, compared to field objects of the same spectral type (from Leggett et al. 2002). Neither of the two TW Hydrae targets harbors measurable excess consistent with an optically thick inner disk.  A large fraction --$\\sim$65\\%-- of brown dwarf candidates in the Trapezium cluster show $K$-band excess (Muench et al. 2001; also see Liu et al., this volume). Our results, albeit for a small sample of objects so far, in the somewhat older $\\sigma$ Ori and TW Hydrae associations suggest lower disk fractions. It may be that small grains in the inner disks have cleared out already by the age of these associations. If so, brown dwarf disks could deplete rapidly, at timescales comparable to or smaller than those for T Tauri disks (Jayawardhana et al. 1999; Haisch et al. 2001). $L^{\\prime}$-band photometry of larger brown dwarf samples in several young clusters, spanning a range of ages, could provide a more definitive answer. Further constraints on disk properties await mid- and far-infrared observations with SIRTF and SOFIA.  Gizis (2002) reported strong H$\\alpha$ emission (equivalent width $\\approx$ 300 \\AA) from one of the TW Hydrae objects, the M8 dwarf 2MASSW J1207334-393254, and suggested it could be due to either accretion, flaring, or chromospheric/coronal activity. Given the lack of substantial $L^{\\prime}$-band excess in this object, accretion now appears less likely as the cause of its strong H$\\alpha$ emission."
    },
    "0206/astro-ph0206279_arXiv.txt": {
        "abstract": "{ In earlier papers (Roshi \\& Anantharamaiah \\cite{ra00}, \\cite{ra01a}), we presented extensive surveys (angular resolution -- 2\\deg\\ $\\times$ 2\\deg\\ \\& 2\\deg\\ $\\times$ 6\\arcmin)  of radio recombination lines (RRLs) near 327 MHz in the longitude range $l = $ 332\\deg\\ $\\rightarrow$ 89\\deg\\ using the Ooty Radio Telescope.  These surveys have detected carbon lines mostly between $l = $ 358\\deg $\\rightarrow$ 20\\deg\\ and in a few positions at other longitudes.  This paper presents the observed carbon line parameters in the high-resolution survey and a study of the galactic distribution and angular extent of the line emission observed in the surveys.  The carbon lines detected in the surveys arise in ``diffuse'' \\CII regions.  The \\lv diagram and radial distribution constructed from our carbon line data shows similarity with that obtained from hydrogen recombination lines at 3 cm from \\HII regions indicating that the distribution of the diffuse \\CII regions in the inner Galaxy resembles the distribution of the star-forming regions.  We estimated the \\FIRCII emission from diffuse \\CII regions and find that upto 95 \\% of the total observed \\FIRCII emission can arise in diffuse \\CII regions if the temperature of the latter $\\sim 80$ K. Our high-resolution survey data shows that the carbon line emitting regions have structures on angular scale $\\sim$ 6\\arcmin.  Analysis of the dual-resolution observations toward a 2\\deg\\ wide field centered at $l = $ 13\\deg.9 and toward the longitude range $l = $ 1\\deg.75 to 6\\deg.75 shows the presence of narrow ($\\Delta V \\le $ 15 \\kms) carbon line emitting regions extending over several degrees in $l$ and $b$.  The physical size perpendicular to the line-of-sight of an individual diffuse \\CII region in these directions is $>$ 200 pc. ",
        "introduction": "Radio recombination lines (RRLs) of hydrogen, helium and carbon have been unambiguously identified in the spectra obtained toward \\HII regions (see review by Roelfsema \\& Goss \\cite{rg92}).  The hydrogen and helium recombination lines mostly originate in hot ($T_e \\sim$ 5000 -- 10000 K) regions ionized by photons of energy $\\ge$ 13.6 eV.  Since the ionization potential of carbon is 11.4 eV, low energy photons (11.4 eV $ \\le E < $ 13.6 eV) that escape from \\HII regions can ionize gas phase carbon atoms outside the hot regions. Thus ionized carbon regions can exist in dense (hydrogen nucleus density $n_0 \\sim$ 10$^{5}$ \\cmthree) photo-dissociation regions (PDRs) adjacent to \\HII regions or in the neutral components (\\HI or molecular) of the interstellar medium (ISM). Tielens and Hollenbach (\\cite{th85}) define PDRs as regions where the heating or/and chemistry of the predominantly neutral gas is governed by the FUV (6--13.6 eV) photons.  Since the FUV photons are omnipresent, the PDRs, by definition encompass a substantial fraction of atomic gas in a galaxy (Hollenbach \\& Tielens \\cite{ht97} and references therein). The dense PDRs (Tielens \\& Hollenbach \\cite{th85}) are located at the interface of molecular clouds and \\HII regions whereas the low-density ($n_0 \\sim 10^3$ \\cmthree) PDRs (Hollenbach, Takahashi \\& Tielens \\cite{htt91}) are located in the diffuse interstellar gas; the ambient FUV flux sufficing to control its chemistry and heating. The ionized carbon regions in the dense PDRs are referred to as ``classical'' \\CII regions. These \\CII regions are observationally identified by the narrow ( 1 -- 10 \\kms) emission lines of carbon at frequencies $>$ 1 GHz. Several studies have been made to understand and model the line emission from such regions (eg. Garay \\etal \\cite{gar98}, Wyrowski \\etal \\cite{wetal00}). These regions are not accessible to low frequency RRLs due to the increased pressure broadening ($ \\propto \\nu^{-8.2/3}$; Shaver \\cite{s75}) and increased free-free continuum optical depths ($\\tau \\propto \\nu^{-2}$). The second class of \\CII regions, referred to as ``diffuse'' \\CII regions, coexists with the diffuse neutral component of the ISM. The emission measures of these regions are fairly low ($< 0.1$ \\cmsix pc;  Kantharia, Anantharamaiah \\& Payne \\cite{kap98}) and hence these regions are observable in low-frequency RRLs of carbon as either absorption lines or emission lines due to stimulated emission from inverted populations. The diffuse \\CII regions, observed in carbon lines at frequencies $ < 1 $ GHz, are the focus of this paper.  The diffuse \\CII region located in the Perseus arm toward the strong radio continuum source, Cas~A has been extensively studied using low frequency recombination lines of carbon.  In fact, most of our knowledge on this class of \\CII regions has come from these observations. Konovalenko \\& Sodin (\\cite{ks80}) were the first to observe a low-frequency (26.3 MHz) absorption line toward Cas A, which was later correctly identified as the 631$\\alpha$ recombination line of carbon by Blake, Crutcher \\& Watson (\\cite{bcw80}).  Since then, several recombination line observations spanning over 14 to 1400 MHz have been made toward this direction (Kantharia \\etal \\cite{kap98} and references therein).  The predicted smooth transition of carbon lines in absorption at frequencies below 115 MHz to lines in emission at frequencies above 200 MHz has been demonstrated toward this direction (Payne, Anantharamaiah \\& Erickson \\cite{pae89}).  The extensive RRL data collected toward Cas~A has been used in modeling the line-forming gas.  The models show that the carbon RRLs originate in small, relatively cool tenuous regions ($T_e$ = 35--75 K, $n_e$ = 0.05--0.1 \\cmthree, size $\\sim 2$ pc; Payne, Anantharamaiah \\& Erickson \\cite{pae94}) of the ISM. Comparison of the distribution of carbon RRLs near 327 MHz observed with the VLA (2\\arcmin.7 $\\times$ 2\\arcmin.4) toward Cas A with \\HI absorption in the same direction suggests that the carbon line-forming region likely coexists with the cold, diffuse \\HI component of the ISM (Anantharamaiah \\etal \\cite{aepk94}).  In addition to the region toward Cas~A, the distribution of the diffuse \\CII regions in the Galaxy has also been studied to some extent.  Surveys have been conducted near 76 MHz ($n \\sim 441$) with the Parkes 64m telescope (Erickson, McConnell \\& Anantharamaiah \\cite{ema95}) and near 35 MHz ($n \\sim 580$) with the Gauribidanur telescope (Kantharia \\& Anantharamaiah \\cite{ka01}) to search for carbon recombination lines, mostly in the inner part of the Galaxy. These observations have succeeded in detecting carbon RRLs in absorption from several directions in the galactic plane with longitudes ranging from $l = $ 340\\deg\\ $\\rightarrow$ 20\\deg. The diffuse \\CII regions appear to be fairly widespread in the inner part of our Galaxy. Observations away from the Galactic plane have shown the region to be several degrees wide in galactic latitude. The positions with detections near 35 MHz were observed near 327 MHz using the Ooty Radio Telescope by Kantharia \\& Anantharamaiah (\\cite{ka01}) and the emission counterparts of the carbon absorption lines were detected. Combining their observations with all other existing carbon RRL observations, they modeled the line emission at different positions in the galactic plane. While models with physical properties similar to those obtained in the direction of Cas A can fit the observed data, the possibility of carbon lines originating in regions with temperature $\\le$ 20 K cannot be ruled out (Kantharia \\& Anantharamaiah \\cite{ka01}).  If the temperature of the diffuse \\CII regions is found to be low, then these regions could even be associated with the molecular component of the ISM (Konovalenko \\cite{k84}, Golynkin \\& Konovalenko \\cite{gk90}, Sorochekov \\cite{s96}, Kantharia \\& Anantharamaiah \\cite{ka01}).  These low temperature regions may be low-density PDRs (Hollenbach, Takahashi, Tielens \\cite{htt91}) formed on surfaces of molecular clouds due to ionization from background FUV radiation. Although some modeling of these diffuse \\CII regions using low-frequency carbon RRLs has been done, a wide range of parameter space has been found to fit the existing observations.  The physical properties, distribution and association of these regions with other components of the ISM requires more investigation. In addition to carbon RRLs, ionized carbon is also traced by the \\FIRCII line. The $158\\mu$ line emission from the Galaxy has been mapped by Bennett \\etal (\\cite{ben94}) and  Nakagawa \\etal (\\cite{naka98}).  They find that the \\FIRCII emission consists of compact emission regions associated with compact \\HII regions (Nakagawa \\etal \\cite{naka98}) and a diffuse emission whose origin is not very clear. Since both the fine-structure line and the carbon RRLs require ionized carbon regions, it is possible that the two can arise from similar regions. Kantharia \\& Anantharamaiah (\\cite{ka01}) tried to compare the carbon lines near 35 MHz with the \\FIRCII emission but they did not derive any conclusive results. Hence, no detailed comparative study of the radio and FIR line emission of carbon from diffuse \\CII regions exists.  In this paper, we also attempt a discussion on these two tracers of ionized carbon regions.  Extensive surveys of recombination lines near 327 MHz have been made with the primary objective to study the low-density ionized gas in the Galaxy by observing low-frequency hydrogen RRLs from this gas (Roshi \\& Anantharamaiah \\cite{ra00}; hereafter Paper I; Roshi \\& Anantharamaiah \\cite{ra01a}; hereafter Paper II; Roshi \\& Anantharamaiah \\cite{ra01b}).  Since the velocity coverage of these surveys was sufficient to allow detection of carbon RRLs, which are separated from the hydrogen line by $\\sim -150$ \\kms\\ , the surveys have succeeded in detecting carbon features toward several positions in the galactic plane.  The surveys have data with two different angular resolutions obtained by using the Ooty Radio Telescope in two different operating modes (see Paper I \\& II). The carbon line data obtained from the higher angular resolution observation (2\\deg\\ $\\times$ 6\\arcmin\\ ) are presented in this paper (see Paper II for spectra) and those obtained in the lower resolution (2\\deg\\ $\\times$ 2\\deg\\ ) survey were presented in Paper I. In this paper, we present a study of the distribution and angular extent of the carbon line forming region in the galactic plane by making use of the carbon RRLs detected in the 327 MHz surveys. Interestingly, in several directions the carbon line emission observed in the surveys seems to be associated with \\HI\\ self-absorption features, which will be discussed in Roshi, Kantharia \\& Anantharamaiah (\\cite{rka02}).  A summary of the observations and basic results are presented in Section~\\ref{sec:obs}. Section~\\ref{sec:dis} discusses the distribution of the diffuse \\CII regions in the galactic disk and compares it with the distribution of other components of the ISM. Section~\\ref{finestructure} discusses the possibility of a common origin of the carbon RRL and the diffuse \\FIRCII line emission. The latitude extent of carbon line emission is discussed in Section~\\ref{sec:lat}.  The higher resolution observations are used to study the angular extent of the carbon line emitting region, which is discussed in Section~\\ref{sec:ang}. Section~\\ref{sec:sum} summarizes the paper. ",
        "conclusions": ""
    },
    "0206/astro-ph0206423_arXiv.txt": {
        "abstract": "s{ The standard afterglow model neglects the presence of the GRB prompt radiation ahead of the blast wave. In fact, the blast wave is dramatically influenced by the leading gamma-ray front which preaccelerates the ambient medium and loads it with electron-positron pairs. The front sweeps the medium outward with a high Lorentz factor and results in a spectacular effect: the GRB ejecta moves freely in a cavity behind its own radiation front. When the front expands sufficiently and gets diluted, a blast wave develops, and it does it differently from the standard model used before. The afterglow should initially appear as a steep rise of soft emission (from infrared to soft X-rays) at a time comparable to the prompt GRB duration and then the emission should quickly evolve to a normal X-ray afterglow. This may explain the prompt optical flash observed in GRB~990123 and allows one to infer the ejecta Lorentz factor in this burst: $\\Gej\\approx 200$. The effect of the gamma-ray front on the afterglow emission is especially pronounced if the GRB has a massive progenitor and the blast wave propagates in the progenitor wind. We emphasize the importance of early afterglow observations in soft bands, as they will allow one to test different progenitor models. } ",
        "introduction": "The afterglow emission of gamma-ray bursts (GRBs) is believed to come from a blast wave driven by a relativistic explosion ejecta into an ambient medium. In contrast, the gamma-ray burst itself is emitted early, probably preceding the development of the blast wave (see [1] for a recent review). The $\\gamma$-rays decouple from the explosion ejecta, overtake it, and form a precursor that interacts with the ambient medium ahead of the blast wave. The leading radiation front and the ejecta are geometrically thin shells, of radius $R$ and thickness $\\Delta\\ll R$, and they are separated by a small distance $l\\approx R/\\Gamma^2\\ll R$ where $\\Gamma$ is the ejecta Lorentz factor. How does this couple of shells interact with the ambient medium?  In the radiation front, two processes take place: Compton scattering and $\\gamma-\\gamma$ absorption of the decollimated scattered photons. As a result the medium is pushed ahead and loaded with $e^\\pm$ pairs~[2-4]. The medium dynamics in the front is nonlinear because the created pairs do more scattering, and the problem is further complicated by the transfer of scattered radiation through the medium. Yet an accurate solution can be obtained~[4] and well approximated by a simple analytical model which is summarized below. It turns out that the radiation front changes the medium faced by the blast wave at radii $R<2\\times 10^{16}E_{53}^{1/2}$~cm (where $E$ is the ``isotropic'' energy of the GRB) and has a dramatic impact on the early afterglow. Most of the afterglow energy is likely emitted inside this radius if the GRB has a massive progenitor. This opens new prospects for testing progenitor models by studying the early afterglow. ",
        "conclusions": "The initial peak of a GRB afterglow is difficult to observe because it happens early. In X-rays, it probably overlaps with the prompt GRB. In soft bands, it is difficult to observe the burst quickly. Future observations may overcome this technical difficulty. {\\em Swift} will be able to see optical emission 50~s after the beginning of the burst, and the proposed {\\em ECLAIRs} can detect flashes at much smaller times. A detected peak of the afterglow provides valuable information on the ejecta Lorentz factor, and its time profile can give indications on the nature of the GRB progenitors.  The $\\gamma$-ray front plays a crucial role for the early afterglow. If the internal scenario of the prompt GRB [1] is correct then the $\\gamma$-rays must clear a cavity in the ambient medium and preclude any afterglow emission at early times. The afterglow should suddenly ``switch on'' and rise sharply at $\\tobs\\sim 10E_{53}^{1/2}(\\Gej/100)^{-2}$~s. The initial flash must be very soft, and it can emit a lot of energy before a normal X-ray afterglow sets in, especially in the massive progenitor scenario.  At the afterglow radii, the ambient medium is optically thin (even after $e^\\pm$ loading), and the bulk of GRB radiation passes freely through it. However, the most energetic $\\gamma$-rays may encounter a substantial $\\gamma-\\gamma$ opacity made by the scattered radiation [4]. In the massive progenitor scenario, prompt $\\gamma$-rays should be absorbed above $\\epsilon_{\\rm br}=5-50$~MeV that depends on the density of the progenitor wind. It produces a spectral break which should be easily detected by {\\em GLAST}. Additional diagnostics is possible if the scattered $\\gamma$-rays are directly observed as an echo of the burst~[9].  The optically thin medium scatters only a small portion of the $\\gamma$-ray energy, much smaller than the energy of the blast wave. In view of this fact, the strong dynamical impact of the $\\gamma$-ray front on the blast wave might seem surprising. The clue to this paradox is the high Lorentz factor of the ejecta, $\\Gej\\gg 1$. In the standard model, a low-mass ambient material efficiently decelerates the relativistic ejecta: half of the explosion energy $\\Eej$ is dissipated when the ejecta sweeps a static $m=\\Eej/\\Gej^2c^2$. If the ejecta sends ahead a radiation precursor, it easily preaccelerates the medium to $\\gamma\\gg 1$ by depositing energy $\\gamma mc^2\\ll\\Eej$. Even at $\\gamma>\\Gej$, the deposited energy is small, while the dynamical impact is enormous: $m$ runs away and the ejecta moves freely in the cleared cavity. Thus dissipation is delayed until the precursor is diluted sufficiently by side expansion."
    },
    "0206/astro-ph0206109_arXiv.txt": {
        "abstract": "{ Effects of rotational mode coupling on photometric parameters of stellar oscillations are studied. At moderate rotation rates, a strong coupling between modes of spherical harmonic degree, $\\ell$, differing by 2 and of the same azimuthal order, $m$, takes place if the frequencies are close. This is a common situation amongst main sequence pulsators. Numerical results for a sequence of $\\beta$ Cephei star models are reported for the two- and three-mode couplings.\\\\ One consequence of mode coupling is that modes of higher degree should be considered in photometric mode identification. Modes with nominal degree $\\ell>2$ acquire substantial $\\ell\\le2$ components and therefore are more likely to reach detectable amplitudes. Coupled mode positions in the amplitude ratio -- phase difference diagrams, based on multicolour photometry, become both aspect- and $m$-dependent. Examples of the mode path in the diagram with varying aspect are given. The diagrams remain a useful tool for mode identification in rotating stars but the tool must be used with care. \\\\ ",
        "introduction": "Mode identification, that is, determination of the radial order and spherical harmonic, is an essential step in asteroseismology. The task is not easy in the case of the oscillation frequency spectra in $\\beta$ Cep and $\\delta$ Sct stars, which are most often lacking equidistant patterns. The photometric diagnostic diagrams, i.e., the amplitude ratio $vs.$ phase difference dependencies in different passbands, are the most popular tools for mode identification in pulsating stars. Following pioneering works (Balona \\& Stobie 1979, Stamford \\& Watson 1981) these tools have been applied mainly to $\\beta$ Cep and $\\delta$ Sct variables.  Theoretical diagnostic diagrams are based on linear nonadiabatic calculations of stellar oscillations and on models of static plane-parallel atmospheres. In the early works an arbitrary parametrization has been used instead of linear nonadiabatic calculations. This approach has been followed even in some recent studies (e.g. Garrido 2000). The nonadiabatic calculations were first included explicitly by Cugier et al. (1994) and subsequently by Balona \\& Evers (1999), Cugier \\& Daszy\\'nska (2001), Balona et al. (2001) and also by Townsend (2002), who applied them to SPB stars.  Up to now, the amplitudes and phases, which we -- following Cugier et al. (1994) -- call photometric nonadiabatic observables, were calculated in the framework of linear nonadiabatic theory, ignoring effects of rotation. However, amongst $\\beta$ Cep and $\\delta$ Sct stars slow rotators, for which such an approximation is adequate, are more an exception than a rule. Here we examine effects of moderate rotation on theoretical diagnostic diagrams. By moderate we mean so slow that perturbational treatment of rotation is adequate. Specifically, we rely on the third order formalism of Soufi et al. (1998).  The most important effect of rotation in the context of diagnostic diagrams is coupling between close frequency modes of spherical harmonic degree , $\\ell$, differing by 2. The effect was discussed in some detail by Soufi et al. (1998). The essential formulae are recalled in the next section of this paper.  Numerical results presented later concern one selected sequence of $\\beta$ Cep models. On a qualitative level the results are applicable to all stars of this type. Our choice of $\\beta$ Cep stars is motivated not only by the abundance of the observational data but also by the fact that we have credible results from linear nonadiabatic calculations. This is not true for $\\delta$ Sct stars where there are serious uncertainties related to the treatment of convection.  Properties of unstable modes in the selected sequence of models are reviewed in Section 3. Also in this section we discuss the occurence of near resonances between two and three modes as well as certain consequences of mode coupling.  In Section 4 we discuss the visibility in various passbands of modes described by a single spherical harmonic over a wide range of $\\ell$. Examples of diagnostic diagrams for coupled modes are given in Section 5. ",
        "conclusions": "We have seen that even at moderate rotation, mode coupling leads to complications in the diagnostic diagrams used for photometric mode identification. Modes of degrees $\\ell>2$, which are often ignored in identification of peaks in oscillation spectra, may acquire substantial low $\\ell$ components, and are more easily detected. Unlike modes described by a single spherical harmonic, the positions of coupled modes in the diagnostic diagrams depend on the aspect and on the azimuthal order. The positions may be quite confusing, for instance, a mode composed of the $\\ell=0$ and $\\ell=2$ components may appear for a range of aspect angles at the $\\ell=1$ position.  All that is not good news for the photometric mode identification procedure. The problems are not confined to a few cases, but occur at the typical rotation rates encountered in $\\beta$ Cep and $\\delta$ Sct stars. Close frequencies of rotationally coupled modes occur over wide ranges of the instability strips. The implication is that we must be careful in using the diagnostic diagrams for inferring the $\\ell$ values. The diagrams remain useful. After all, they do provide observational constrains on stars and their oscillations. However, mode identification may be done only simultaneously with determination of stellar parameters and inclination of rotational axis."
    },
    "0206/nucl-th0206051_arXiv.txt": {
        "abstract": "{ We solve analytically the ellipsoidally expanding fireball hydrodynamics with source terms in the momentum and energy equations, using the non-relativistic approximation. We find that energy transport from high $p_t$ jets of gluons to the medium leads to a transient, exponential inflation of the fireballs created in high energy heavy ion collisions. In this transient, inflatory period, the slopes of the single particle spectra are exponentially increasing, while the HBT radius parameters are exponentially decreasing with time. This effect is shown to be  similar to the development of the homogeneity of our Universe due to an inflatory period. Independently of the initial conditions, and the exact value of freeze-out time and temperature, the measurables (single particle spectra, the correlation functions, slope parameters, elliptic flow, HBT radii and cross terms) become time-independent during the late, non-inflatory stages of the expansion, and they satisfy new kind of scaling laws. If the expansion starts with a transient inflation caused by the gluon wind, it leads naturally to large transverse flows as well as to the the simultaneous equality, and scaling behaviour of the HBT radius parameters, $R_{side}\\approx R_{out} \\approx  R_{long} \\approx t_f \\sqrt{T_f/m}$. With certain relativistic corrections, the scaling limit is $\\tau_f \\sqrt{T_f/m_t}$, where $m_t$ is the mean transverse mass of the pair. }   \\keyword{ inflation, fireball hydrodynamics, relativistic heavy ion collisions, single-particle spectra, Bose-Einstein correlations } \\PACS{ 12.38.Mh; 24.85.+p; 25.75.-q, ... }  \\begin{document} ",
        "introduction": "Recently, there were a number of QCD based perturbative calculations that indicated the energy loss of gluon jets with high transverse momentum due to multiple scattering during the course of their penetration through hot and dense hadronic matter. This phenomena, the jet quenching goes back to papers of Gyulassy, Pl\\\"umer and Wang~\\cite{Wang:1994fx,Gyulassy:1993hr}, recently studied in the papers of Gyulassy, L\\'evai  and Vitev (GLV)~\\cite{Gyulassy:2000fs}. It is clear from these works, that the jet quenching mechanism results in the energy loss or the depletion of the number of hadrons from the high transverse momentum part of the spectrum. Due to the conservation of the total energy, this implies that the number of particles with small total momentum has to increase. See ref. \\cite{Vitev:2002vr} for the summary of recent developments in this field from the point of view of perturbative QCD.  In this paper, we focus on the phenomenological consequences of energy and momentum transfer to the soft, hydrodynamically behaving, small transverse momentum part of the single particle spectrum. We present a new family of solutions of non-relativistic hydrodynamics with source terms. This family of solutions is a generalization of the results of refs.~\\cite{ellobs,ellsol,cssol} to ellipsoidally symmetric, expanding fireballs that are subject to energy and momentum transfer from the non-hydrodynamically behaving modes. A constant rate of energy pumping is shown to lead to an exponential inflation of the principal axis of the fireballs, while momentum transfer leads only to linear increase. These effects are correlated in time to the transition of high energy gluons through the nuclear medium, hence, by default, they are also transient phenomena, existing only in the initial phase of the expansion. We point out the similarity of this mechanism to the existence of a transient inflatory period in the expansion of the early Universe.  Inflation of the Universe after the Big Bang leads to a homogeneous and flat solution of Einstein's equation. We show here that an inflation of hadronic fireballs, caused by the gluon wind in the initial stage of the Little Bangs of heavy ion collisions also leads to the simultaneous equality, spherical symmetry and scaling of the effective, measurable source sizes (frequently referred to as the HBT radii). This effect connects the physics of the smallest observable scales to the physics of the largest observable ones, underlining the universal nature of the laws of physics and the scale independence of the local conservation laws, governing the hydrodynamics of these expansions. ",
        "conclusions": "We find that an inflatory period automatically leads to a spherical symmetry for the observable HBT radius parameters in an analytically solvable model of fireball hydrodynamics with source terms motivated by the flux of high-pt gluon wind through the expanding hadronic matter, created in high energy collisions of heavy ions. The effect is very similar in nature how a flat and spatially homogeneous Universe is obtained after a period of inflation in astrophysics. Our results underline the similarity between the physics of 'Little Bangs' at the smallest experimentally accessible scales and the physics of the 'Big Bang' of our Universe, at the largest observable scales.  We also find, that the form of the observables is {\\it exactly} the same for cases with or without inflation, and for various changes of the equation of state in the intermediate steps of the time evolution.  This implies that various different equations of state from various different initial conditions may lead to {\\it exactly} the same hadronic final state. Furthermore, the time evolution of the parameters of the hydrodynamical solution is sensitive not only to the equations of state but also to the magnitude of the source terms in the energy and momentum balance equations, which may provide non-equilibrium mechanisms to connect the initial state to a violently exploding final state.  These results imply that {\\it i)} the hadronic final state of fireball evolution does not remember the path (hence any earlier transient phase of matter) of its time evolution, but {\\it ii)} given the equation of state, and given the possible source terms in the energy and the Euler equation, the initial state can still be reconstructed, from the observables in the final hadronic state.  Further studies are in progress in for the publication of the ellipsoidally symmetric solutions of relativistic hydrodynamics and for the calculation of the observables corresponding to these relativistic generalizations of the present results.   {\\it Acknowledgments:} This study has been supported by the OTKA grants T034262, T038406,  by the US - Hungarian NSF - OTKA - MTA grant 0089462 and by an Alumni Initiatives Award of the Fulbright Foundation."
    },
    "0206/hep-ph0206266_arXiv.txt": {
        "abstract": "{Coannihilation processes provide an important additional mechanism for reducing the density of stable relics in the Universe. In the case of the stable lightest neutralino of the MSSM, and in particular the Constrained MSSM (CMSSM), the coannihilation with sleptons plays a major role in opening up otherwise cosmologically excluded ranges of supersymmetric parameters. In this paper, we derive a full set of exact, analytic expressions for the coannihilation of the lightest neutralino with the sleptons into all two--body tree--level final states in the framework of minimal supersymmetry. We make no simplifying assumptions about the neutralino nor about sfermion masses and mixings other than the absence of explicit CP--violating terms and inter--family mixings. The expressions should be particularly useful in computing the neutralino WIMP relic abundance without the approximation of partial wave expansion. We illustrate the effect of our analytic results with numerical examples and demonstrate a sizeable difference with approximate expressions available in the literature.  } ",
        "introduction": "The relic density of stable weakly--interacting massive particles (WIMPs) is determined primarily by how efficiently their number density  in the early Universe can be reduced. In the case of the most popular WIMP candidate: the lightest neutralino of minimal supersymmetry (SUSY), assumed to be the lightest SUSY particle (LSP), there are two generic mechanisms~\\cite{kt90,jkg96}. First, the neutralino can pair--annihilate into ordinary particles. Second, in some cases they can coannihilate with some some other species if these are nearly mass--degenerate with the LSPs.  The standard mechanism of neutralino pair--annihilation has been considered in much detail in many papers~\\cite{jkg96}. In particular, complete sets of neutralino annihilation cross sections were provided in~\\cite{dn93} (see also~\\cite{jkg96}) in the case of partial wave approximation. Exact expressions applicable both near resonances and new final--state thresholds were recently published in~\\cite{nrr2}.  The mechanism of coannihilation was originally pointed out by Griest and Seckel~\\cite{gs91}. It applies when there exists some other species $\\chi^\\prime$ which is not much heavier than the WIMP and may therefore be still present in the thermal plasma at the time of WIMP decoupling. Coannihilation becomes important if its annihilation with the WIMP (and/or itself) is equally, or more, efficient than the pair--annihilation of the stable WIMPs.  These circumstances in particular are naturally realized in the case of the higgsino--like lightest neutralino LSP in the framework of the Minimal Supersymmetric Standard Model (MSSM).  In this case the next--to--lightest neutralino and the lightest chargino are almost mass--degenerate with the LSP~\\cite{mizuta92,eg97}. In fact, the coannihilation in this case is so efficient that it has a devastating effect on the relic density of the higgsino--like neutralino, which is the type of LSP strongly disfavored by naturalness~\\cite{chiasdm} and mass--unification~\\cite{rr93,an93,kkrw94}.  It even plays some role~\\cite{eg97} in the strongly prefered case of the bino--like LSP~\\cite{chiasdm}.  Since, in the framework of general softly--broken low--energy SUSY, scalar superpartner masses are a priori unrelated (or at best loosely related) to the neutralino sector, in principle one might assume that any of the scalar superpartners could be light enough to participate in coannihilation with the neutralino LSP -- the case that would be technically rather challenging and in any case not particularly well--motivated.  A more realistic scenario is the one in which one of the scalar partners of the top--quark or the $\\tau$--lepton is rather light and nearly degenerate in mass with the neutralino LSP. This is because the off--diagonal elements in their $2\\times2$ mass matrices can under some circumstances greatly reduce one of the eigenmasses relative to the other.  The case of neutralino--stop coannihilation was considered in~\\cite{bdd00} in the framework of the MSSM and recently shown in~\\cite{eos01} to be also applicable in the case of the Constrained MSSM (CMSSM) but only for rather large values of the trilinear soft SUSY--breaking term $\\azero$.  The importance of the neutralino coannihilation with the lighter of the two staus in the framework of the CMSSM was pointed out in~\\cite{efo98}. In this model, when the common gaugino mass parameter $\\mhalf$ is much larger than the common scalar mass parameter $\\mzero$, it is the lightest stau that is the LSP~\\cite{kkrw94}. The effect of the neutralino--stau coannihilation is to open up a narrow corridor~\\cite{efo98} just above the boundary of equal neutralino--stau masses into an otherwise cosmologically forbidden region in the $(\\mhalf,\\mzero)$--plane.  (Without coannihilation, the requirement of the relic abundance of the neutralino to be consistent with that allowed by the age of the Universe ($\\abundchi\\lsim{\\cal O}(1)$) often provides a stringent upper bound on the parameters $\\mhalf$ and $\\mzero$~\\cite{rr93,kkrw94}.) The effect has since been included in a number of recent analyses, \\eg\\ in~\\cite{coann:recent,bbb02}, and in a publically available package for computing the relic density micrOMEGAs~\\cite{micromegas}.  One should mention that there are two other ways of evading this otherwise generic cosmological bound on $\\mhalf$ and $\\mzero$. One is realized for $\\mzero\\gg\\mhalf$.  In this region one invariably finds it difficult to satisfy the conditions of radiative electroweak symmetry breaking (EWSB); in other words the square of the Higgs/higgsino mass parameter $\\mu$ comes out to be negative. In a very narrow corridor along the region of no--EWSB, $\\mu$ grows rapidly from zero but it is still less than $\\mhalf$~\\cite{kkrw94,fmw00}. As a result, the LSP has a sizeable higgsino component (although it is still mostly a bino, like in the rest of the $(\\mhalf,\\mzero)$--plane) and its relic density is typically small. In fact, because of the growing LSP mass and its gaugino fraction, the relic density increases rapidly along a very steep slope from very small values, characteristic for lighther neutralinos with a larger higgsino admixture, to larger (and often cosmologically excluded) values characteristic of heavier and bino--dominated neutralino. As a result, the cosmologically favored range $0.1<\\abundchi<0.2$ is only realized there for a rather narrow range of $\\mhalf$~\\cite{fmw00,rrn1}.  The other important escape route from the cosmological bound on $\\mhalf$ and $\\mzero$ occurs when $\\tanb$, which is the usual ratio of the vacuum expectation values of the neutral Higgs scalars, is large, $\\gsim 50$.  This is because the physical masses of the heavy Higgs scalar $\\hh$ and pseudoscalar $\\ha$ decrease with increasing $\\tanb$~\\cite{dn93}.  When the neutralino mass becomes large enough, close to half of the heavy Higgs boson mass, the LSP relic abundance becomes efficiently reduced through a relatively wide resonance involving mostly the pseudoscalar exchange. The effect is amplified by the coupling of $A$ to down--type fermions which grows like $\\tanb$. (The heavy scalar Higgs coupling also grows in a similar fashion but its dominant contribution is only $p$--wave and therefore suppressed by square of the WIMP relative velocity.) In the framework of the CMSSM, the effect is to open up wide fractions of the otherwise cosmologically excluded ranges of the $(\\mhalf,\\mzero)$--plane along the wide $A$--resonance~\\cite{efo98,rrn1,bk01,ls01,ddk01}. The precise position of the resonance shows a sizeable dependence on some input parameters, most notably on the ratio $\\mt/\\mb$, $\\azero$, \\etc, and remains a subject of some ongoing debate. The full two--loop Higgs effective potential would have to be computed and implemented in the analysis to reduce the sensitivity to, \\eg, the scale dependence.  In the CMSSM a set of reasonably well--motivated unification assumptions leads to only four parameters: a common gaugino mass $\\mhalf$, a common scalar mass $\\mzero$, a trilinear coupling $\\azero$, as well as $\\tanb$. One is also free to choose the sign of the $\\mu$--parameter, while its magnitude is determined by the mechanism of electroweak radiative symmetry breaking (EWSB). The CMSSM can therefore be considered as a well--motivated SUSY model with the smallest number of independent parameters. (In particular, the so--called minimal supergravity (mSUGRA) model can be viewed as a specific realization of the framework.)  The CMSSM has become a benchmark model for the LHC and other SUSY searches.  As mentioned above, the neutralino--stau coannihilation effect is particularly important in the framework of the CMSSM, but it can affect the neutralino relic density also in the more general MSSM. This will be the framework in which we will work here for the sake of generality.  In this paper, we will present a full set of exact, analytic expressions for the tree--level cross section of the neutralino coannihilation with sleptons into all two--body final states in the general MSSM. In our analysis we will make no simplifying assumptions about the neutralino, nor will we assume the degeneracy of the left-- and right--sfermion masses. We will not consider here the possibility of CP and flavor violation in the slepton sector although we will assume a general form of the left--right slepton mixing within each generation.  We will include all tree--level final states and all intermediate states. We will also keep finite widths in $s$--channel resonances.  A set of expressions for the neutralino--slepton coannihilaton was given in~\\cite{efo98,efos00} but only in the approximation of the partial wave expansion. Furthermore, these formulae did not include the effects of the tau Yukawa, of the $\\stauone-\\stautwo$ mixing and in some channels of the mass of the $\\tau$, \\etc, which make them less reliable at large $\\tanb\\gsim20$~\\cite{efgos01}.  The results presented here are exact, include all the terms and are valid both small and large values of $\\tanb$, and both near and away from resonances and thresholds for new final states. This paper is meant to be a follow-up to~\\cite{nrr2} where we have calculated all the analytic cross sections of all tree--level processes for the neutralino pair--annihilation into all two--body final states. We follow the same conventions and notations as in~\\cite{nrr2}.  The plan of the paper is as follows. In Sect.~\\ref{relicdensity:sec} we briefly review the formalism for computing the relic density in the presence of both annihilation and coannihilation.  In Sect.~\\ref{mssm:sec} we introduce the relevant ingredients of the MSSM and list all the neutralino pair--annihilation channels. Explicit expressions for the coannihilation cross secion are given in Sect.~\\ref{exact:sec}. In Sect.~\\ref{numanalysis:sec} we present some numerical examples and in Sect.~\\ref{summary:sec} we summarize our work. Appendix~A contains a complete list of relevant couplings and in Appendix~B we provide expressions for several auxiliary functions used in the text. ",
        "conclusions": "\\label{summary:sec} The accuracy of determining the abundance of the dark matter in the Universe is continuously improving. This requires theoretical computations of the neutralino relic abundance to be performed with at least the same, if not better, level of precision, if one wants to reliably compare theoretical predictions with observations.  In this paper we have derived a full set of exact, analytic expressions for the neutralino--slepton coannihilation cross sections into all tree--level two--body final states. While these formulae are applicable in the framework of the general MSSM, they are of particular importance in the context of the Constrained MSSM. In this framework, which is often considered a ``reference'' SUSY model and is thus of much interest to the community, much of the allowed regions are a result of the neutralino--slepton coannihilation. Our results should help in allowing one to determine these regions more precisely.   \\bigskip"
    },
    "0206/astro-ph0206364_arXiv.txt": {
        "abstract": "{Based on the Bica et al.\\ (\\cite{bica}) catalogue, we studied the star cluster system of the LMC and provide a new catalogue of all binary and multiple cluster candidates found. As a selection criterion we used a maximum separation of $1\\farcm4$ corresponding to 20 pc (assuming a distance modulus of 18.5 mag). We performed Monte Carlo simulations and produced  artificial cluster distributions that we compared with the real one in order to check how many of the found cluster pairs and groups can be expected statistically due to chance superposition on the plane of the sky. We found that, depending on the cluster density, between $56 \\%$ (bar region) and $12\\%$ (outer LMC) of the detected pairs can be explained statistically. We studied in detail the properties of the multiple cluster candidates. The binary cluster candidates seem to show a tendency to form with components of similar size. When possible, we studied the age structure of the cluster groups and found that the multiple clusters are predominantly young with only a few cluster groups older than 300 Myr. The spatial distribution of the cluster pairs and groups coincides with the distribution of clusters in general; however, old groups or groups with large internal age differences are mainly located in the densely populated bar region. Thus, they can easily be explained as chance superpositions. Our findings show that a formation scenario through tidal capture is not only unlikely due to the low probability of close encounters of star clusters, and thus the even lower probability of tidal capture, but the few groups with large internal age differences can easily be explained with projection effects. We favour a formation scenario as suggested by Fujimoto \\& Kumai (\\cite{fk}) in which the components of a binary cluster formed together and thus should be coeval or have small age differences compatible with cluster formation time scales. ",
        "introduction": "\\label{intro_montecarlo}  The first systematic work on binary clusters in the Magellanic Clouds started approximately a decade ago. The first catalogue of binary star clusters in the LMC was presented by Bhatia \\& Hatzidimitriou (\\cite{bh}) and Bhatia et al.\\ (\\cite{brht}) who surveyed the cluster system (consisting of 1200 objects known at that time) and listed 69 binary cluster candidates. Their selection criterion was a maximum separation between the components of a proposed pair of approximately 18 pc (assuming a distance modulus of 18.4 mag). Following Page (\\cite{page}), out of these 69 double objects only 31 can be explained as optical pairs, i.e., clusters that appear as close pairs on the plane of the sky due to projection effects. Ages were available only for some of the clusters and suggested that the pairs are young (between $10^{7}$ to a few $10^{8}$ yr), consistent with expected time scales for merger or disruption of binary clusters (Bhatia \\cite{bhatia}).  In the following years, more studies on binary cluster candidates were performed but concentrated mainly on one or a few individual objects in order to establish their binarity (Kontizas et al.\\ \\cite{kkx}, Lee \\cite{lee92}, Bhatia \\cite{bhatia92}, Kontizas et al.\\ \\cite{kkm}, Vallenari et al.\\ \\cite{vafcom}, Hilker et al.\\ \\cite{hrs}, Grebel \\cite{gr97}, Vallenari et al.\\ \\cite{vbc}, Leon et al.\\ \\cite{lbv}, Dieball \\& Grebel \\cite{dg, dg2000}, Dieball et al.\\ \\cite{dgt}). Few theoretical studies concerning formation, gravitational interaction and dynamical evolution of binary clusters are available (Sugimoto \\& Makino \\cite{sm}, Bhatia \\cite{bhatia}, Fujimoto \\& Kumai \\cite{fk}, de\\,Oliveira et al.\\ \\cite{dodb}, Theis \\cite{theis}).  However, since the investigation of Bhatia \\& Hatzidimitriou (\\cite{bh}) and Bhatia et al.\\ (\\cite{brht}) many more clusters have been discovered in the LMC. Thus, it is time to perform a new study on the nowadays better known LMC cluster system, aiming at the question of how many close cluster pairs exist and how many of these might be explained as chance superpositions.  Recently, Pietrzy\\'{n}ski \\& Udalski (\\cite{pu}) provided a new but spatially limited catalogue of multiple cluster candidates in the LMC. They based their studies on the OGLE (see Udalski et al.\\ \\cite{ogle}) dataset which covers 5.8 square degrees of the inner part of the LMC and contains 745 star clusters (Pietrzy\\'{n}ski et al.\\ \\cite{puk}). Out of these, a total of 100 multiple cluster candidates with a maximum separation of 18 pc, assuming a distance modulus of 18.24 mag, were selected. The cluster groups consisted of 73 pairs, 18 triple systems, 5 systems containing four components, 1 with five and 3 systems with six clusters. Assuming that all 745 clusters are distributed uniformly in the 5.8 square degree region and adopting the same statistical approach as Bhatia \\& Hatzidimitriou (\\cite{bh}), 51 chance pairs can be expected. A more detailed investigation of the cluster distribution led to nearly the same result of 53 random pairs. The number of all detected candidates is 153 and thus significantly larger than expected from chance superposition. Ages for the components were taken from Pietrzy\\'{n}ski \\& Udalski (\\cite{pu_lmc}). 102 components are coeval, 53 have very different ages, and most objects are younger than 300 Myr with a peak at 100 Myr. This suggests that most of the multiple clusters have a common origin and are quite young objects.  A catalogue of multiple cluster candidates in the SMC was published by Pietrzy\\'{n}ski \\& Udalski (\\cite{pu_smc}), containing 23 binary and 4 triple cluster candidates. A comparison of both the LMC and SMC binary cluster lists reveals that the distribution of the components' separation, the fraction of cluster groups ($\\approx$ 12 \\%) and their ages are very similar. The similar ages of the binary cluster candidates in both the LMC and SMC might be connected with the last close encounter between these two galaxies. De\\,Oliveira et al.\\ (\\cite{odbd}) presented an isophotal atlas of 75 binary and multiple clusters (comprising 176 objects) from the Bica \\& Dutra (\\cite{bd}) catalogue of SMC clusters. Bica \\& Dutra (\\cite{bd}) included also new discoveries from the OGLE catalogue of SMC clusters (Pietrzy\\'{n}ski et al.\\ \\cite{puk_smc}). Investigating the isophotes of the binary and multiple cluster candidates, de\\,Oliveira et al.\\ (\\cite{odbd}) found isophotal distortions, connecting bridges, or common isophotal envelopes for 25 \\% of the suggested multiple clusters. The authors interpreted this as signs of interaction between the components of a supposed binary or multiple cluster, in agreement with the findings from previous N-body simulations (de\\,Oliveira et al.\\ \\cite{obd}). Ages for 91 out of the 176 clusters that are part of pairs or groups were investigated based on the OGLE $BVI$ maps. 40 clusters are in common with Pietrzy\\'{n}ski \\& Udalski (\\cite{pu_smc}), and de\\,Oliveira et al.\\ (\\cite{odbd}) found good agreement with the study of the OGLE group. Most clusters are young, and the age distribution shows a relevant peak around  200 Myr that can be attributed to the last close encounter between SMC and LMC. The components of groups with more than two members are younger than 100 Myr, which might be an indication that multiple clusters coalesce into binary or single clusters within this timescale. 55 \\% of the binary and multiple cluster candidates were found to be coeval. From this the authors concluded that tidal capture is a rare phenomenon.  In this paper, we present a statistical study of close pairs and multiple clusters in the LMC. We decided to base our analysis on the new, extended catalogue of stellar clusters, associations, and emission nebulae in the LMC provided by Bica et al.\\ (\\cite{bica}, hereafter BSDO). The authors surveyed the ESO/SERC R and J Sky Survey Schmidt films, checked the entries of previous catalogues and searched for new objects. The resolution of the measurements was $< 4\\arcsec$ (Bica \\& Schmitt \\cite{bs}). The resulting new catalogue unifies previous surveys and contains 6659 entries, out of which 3246 are new discoveries that are not mentioned in previous catalogues and lists. Thus, the BSDO catalogue can be considered as the so far most complete catalogue of LMC stellar clusters and associations. We restricted our study to bound stellar systems, which means that we selected only objects which are categorized as ``C''-type (cluster-type), and left out associations and emission nebulae, which are not of interest in the context of the present study. This reduces the number of objects found in the BSDO catalogue from a total of 6659 to 4089.  Based on this catalogue, we performed a statistical study of cluster pairs and groups and provide a complete list of all multiple cluster candidates in the LMC.  In the following sections, we address a number of questions: How many cluster pairs can be found with a projected separation of less than 20 pc between the components of a pair (Sect.~\\ref{groups_montecarlo})? Following Bhatia \\& Hatzidimitriou (\\cite{bh}) and Bhatia et al.\\ (\\cite{brht}), we consider this to be a good selection criterion. Several cluster pairs may form a larger cluster group, e.g., if a component of a cluster pair is less than 20 pc distant from any component of another pair. In this way, three clusters may form a triple cluster, but they also might constitute three cluster pairs if each cluster is seen within 20 pc from each other cluster. How many ``multiple'' clusters, consisting out of more than two single objects, are present, and how many single clusters are involved in these pairs and groups (Sect.~\\ref{groups_montecarlo})? How many of these pairs and multiple systems can be expected statistically, and of how many individual components do they consist (Sect.~\\ref{sim_montecarlo})? Are there any correlations between the properties of the cluster systems such as ages, radii and separations between the components (Sect.~\\ref{groupsproperties})? What is the fraction of coeval pairs or groups compared with the number of multiple clusters whose internal age differences exceed the protocluster survival time (Sect.~\\ref{pairsages})? Does the percentage of coeval systems agree with the number of statistically expected groups (Sect.~\\ref{pairsages})? And finally, do our results favour or give hints at a specific cluster formation scenario (Sect.~\\ref{summary_montecarlo})? For instance, can cluster pairs be explained with statistically expected cluster encounters in the LMC, which lead to tidal capture and thus to bound pairs of different ages (see Sect.~\\ref{encounter_montecarlo})? Or are the multiple cluster candidates predominantly found to be coeval, favouring the formation scenario of Fujimoto \\& Kumai (\\cite{fk}) or of Theis (\\cite{theis})? ",
        "conclusions": "\\label{summary_montecarlo}  We investigated the BSDO catalogue and provide a new catalogue of all binary and multiple cluster candidates found in the LMC. The catalogue is presented in Table~\\ref{groupscatalogue}. Age information available in the literature is also given. We found in total 473 multiple cluster candidates. The separations between the clusters' centres are $\\le 1\\farcm4$ corresponding to 20 pc (assuming a distance modulus of 18.5 mag).  We performed a statistical study of cluster pairs and groups. For this purpose we distinguished between regions of different cluster densities in the LMC.  Vallenari et al.\\ (\\cite{vbc}) and Leon et al.\\ (\\cite{lbv}) proposed that the encounter rate in large cluster groups is higher so that binary clusters can be formed through tidal capture. Such a scenario might explain large age differences between cluster pair components. For each selected region we calculated the encounter rate of star clusters. However, we found that the probabilities for cluster encounters are universally very low. In addition, the probability of {\\it tidal capture} depends on further constraints which will not be fulfilled during every encounter. Thus we conclude that it seems unlikely that a significant number of young pairs may have formed in such a scenario.  We counted the number of all cluster pairs and groups found in the selected areas. In order to check how many of these multiple cluster candidates can be expected statistically due to chance line-up, we performed Monte Carlo experiments for each region to produce artificial cluster distributions which are compared with the real LMC cluster distribution. For all selected regions, the number of chance pairs in our simulations is much lower than the quantity of cluster pairs found: Between $56 \\%$ (in the bar region) and $12 \\%$ (in the outer LMC ring) of all detected pairs can be explained statistically. Especially large cluster groups with more than four members hardly occur in the artificial cluster distributions. A significant number of the cluster pairs and groups cannot be explained with chance superposition and thus might represent ``true'' binary and multiple clusters in the sense of common origin and/or physical interaction.  We studied the properties of the multiple cluster candidates:  In the distribution of the centre-to-centre separations of the cluster pairs two peaks around 6 pc and 15 pc are apparent. This bimodal distribution is more apparent for cluster pairs in which both components have diameters smaller than 7 pc, but cannot be neglected for pairs consisting of larger clusters. We cannot confirm a uniform distribution of separations for pairs with large clusters as suggested by Bhatia et al.\\ (\\cite{brht}). Around separations of 9 -- 10 pc, the number of cluster pairs is depleted. This dip might be interpreted as a balance between the effects that lead to an increase in the number of cluster pairs towards either smaller (due to projection effects) or larger separations (pairs with larger separations are more easily detected).  The size distribution of the group components shows a peak at $0\\farcm45$ ($\\approx 6.6$ pc). Most clusters involved in pairs or groups are small and only few clusters have diameters larger than $1\\farcm8$ (26 pc). The size distribution for group components is very similar to the size distribution for all LMC clusters. It seems that binary clusters tend to form with components of similar size.  The spatial distribution of the multiple cluster candidates coincides with the distribution of clusters in general.  Only for a fraction ($\\approx 27 \\%$) of the clusters that form binary and multiple cluster candidates age information is available, and for only 96 groups ages are known for more than one cluster so that the age structure of the specific group can be examined. For 57 groups the members appear to be either coeval or have ages similar enough to agree, within the errors of the age determination, with a common formation in the same GMC, i.e., the age differences are 10 Myr at maximum (Fukui et al.\\ \\cite{fukui}, Yamaguchi et al.\\ \\cite{yama}). The remaining 39 groups have internal age differences which make a common origin of the components unlikely.  The clusters involved in pairs or groups are found to be predominantly young. The age distribution shows peaks at 4 Myr, 25 Myr and 100 Myr. Our findings differ from Pietrzy\\'{n}ski \\& Udalski (\\cite{pu}) in a way that the two peaks at the younger ages are missing in their age distribution. This is due to the fact that these authors investigated only a part of the LMC and used also a smaller distance modulus that leads to higher ages in general.  We scrambled the ages of the groups components and then randomly assigned them to the group members. On average, $12.9\\pm2.7$ groups with internal age differences $\\le 10$ Myr can be expected, however, 46 groups with internal age differences $\\le 10$ Myr can be found in the real distribution (note that the borderline cases are not considered in this number, see Sect.~\\ref{pairsages}), a number significantly larger than the expected one. Also, the group age distribution for scrambled member ages is smoother than the real one, and the internal age scatter is significantly larger for the groups with random member ages.  No correlation was found between the groups' ages and their internal mean separation. However, there might be a weak tendency towards larger internal age scatter with larger internal separations (indicating larger groups) but a strong tendency as suggested by Efremov \\& Elmegreen (\\cite{ee}) cannot be confirmed.  Most multiple cluster candidates are found to be younger than 300 Myr. A larger number of old cluster groups or of groups with different ages for the components are not found. A formation scenario through tidal capture is not only unlikely due to the very low probability of tidal capture (even in the dense bar region), but the few old groups and the groups with large internal age differences can easily be explained with projection effects, especially since the majority of these groups are located in the dense bar region. Thus, we do not see evidence for an ``overmerging problem'' as proposed by Leon et al.\\ (\\cite{lbv}).  Our findings are clearly in favour of the formation scenario proposed by Fujimoto \\& Kumai (\\cite{fk}) who suggested that the components of a binary cluster formed together, and thus should be coeval or at least have a small age difference compatible with cluster formation time scales."
    },
    "0206/astro-ph0206152_arXiv.txt": {
        "abstract": "I examine the hypothesis that many of the high velocity white dwarfs observed by Oppenheimer et al. (2001) are the remnants of donor stars from binaries that produced type~Ia supernovae via the `single degenerate' channel. If this channel is a significant contributor to the Galactic SN~Ia supernova rate, then the local density of such remnants with $V_{\\perp}>100 \\,{\\rm km.s^{-1}}$ could be as high as $2 \\times 10^{-4} {\\rm pc^{-3}}$, comparable to the densities found by Oppenheimer et al. This white dwarf population differs from others in that it is composed exclusively of single stars. ",
        "introduction": "The notion that white dwarfs may play a part in the Galactic dark matter has been actively investigated in the last few years, motivated primarily by the microlensing observations (Alcock et al. 1997, 2000). Recently there have been claims that this population has been detected directly (Oppenheimer et al. 2001). However, neither of these claims is without detractors. The microlensing optical depth to the LMC may also receive a significant contribution from LMC-associated populations as well (Sahu 1994; Zhao \\& Evans 2000; Alcock et al. 2001), while the high velocity white dwarf sample may possess a thick disk component (Reid, Sahu \\& Hawley 2001; Reyle, Robin \\& Creze  2001; Koopmans \\& Blandford 2001; Hansen 2001).  In this paper I wish to investigate a potential source of high velocity white dwarfs different from the usual population; one which may appear in the searches for halo white dwarfs. The proposed population results from the formation of white dwarfs from stars in the range $1.3-3\\, \\rm M_{\\odot}$ which act as the donor stars in close mass-transfer binaries in which the accretor is a white dwarf. Observationally, these binaries are thought to be the origin of the Super-Soft X-Ray sources. These binaries may also contribute to the type~Ia supernova rate in spiral galaxies. If so, the explosion of the accretor will release the donor star with its pre-supernova orbital velocity. Thus, the white dwarf remnants of these donors (hereafter DR, or `donor remnants') will be a population of high velocity white dwarfs. It is the origin and implications of this population that we investigate below. ",
        "conclusions": "If the majority of Type~Ia supernovae in the Galaxy are the result of single degenerate binaries with slightly evolved main sequence donors, the post supernova kinematics of the donors leads to a population of high velocity white dwarfs. These can potentially explain the interesting observations of Oppenheimer et al. (2001). Not only is the density sufficient to explain the observations, but the age distribution is more characteristic of the thin disk than the thick disk, as is suggested by the Oppenheimer et al. luminosity function (Hansen 2001).  The observational test of this proposal is to search for evidence of binarity amongst the high velocity white dwarfs, by obtaining accurate trigonometric parallaxes and searching for common proper motion main sequence companions, particularly amongst the hotter white dwarfs, where the contribution of the spheroid or thick disk is smaller. Although the DR represent a new population of high velocity white dwarfs and might be considered interesting from the point of view of microlensing, the total mass involved is quite small and the expected contribution to the microlensing optical depth from this population is negligible unless the Galaxy retains only a small fraction of the total iron produced by Type~Ia supernovae."
    },
    "0206/astro-ph0206478_arXiv.txt": {
        "abstract": "The recent detection of resolved radio emission from  AG Dra by MERLIN reported by Ogley et al. is discussed in the context of the wind environment and the physical parameters and geometry of this symbiotic binary system. In particular, it is shown that the two radio components are closely aligned with the binary axis, and their separation suggests their origin in jets ejected from AG Dra during the recent 1995--98 series of oubtbursts. ",
        "introduction": "AG Dra is the well-studied symbiotic binary consisting of a high-velocity, metal-poor, bright K-type giant and a hot white dwarf companion (e.g. Miko{\\l}ajewska et al. 1995; Smith et al. 1996). The binary system has an orbital period of 549 days and a well-defined spectroscopic orbit (Fekel et al. 2001, and references therein). AG Dra is also among the most active symbiotic stars. Its optical light curve is characterized by a series of active (outbursts) and quiescent phases (e.g. Fig. 7 of G{\\'a}lis et al. 1999). Although the activity of AG Dra, and other classical symbiotic stars is still poorly understood, multifrequency observations covering a few whole activity cycles indicate that in AG Dra this activity is related to changes of both radius and temperature of the hot component (e.g. Miko{\\l}ajewska et al. 1995; Greiner et al. 1997). AG Dra is also one of 10 known galactic supersoft X-ray sources (Greiner et al. 1997).  Recently, the star has been searched, together with all the northern supersoft X-ray sources, for radio emission at 5 and 8.4 GHz (Ogley et al. 2002, hereafer O02). The observations by MERLIN telescope have confirmed previous VLA detections. Moreover, the source has been resolved at the milliarcsec scale into two components of nearly equal brightness, and combined flux of $\\sim 1$ mJy. O02 have also studied possible interpretations of this emission in terms of a wind environment from either the cool giant or the hot white dwarf. They have concluded that all their scenarios give the radio emission fluxes an order of amplitude lower than the observed value.  In the following we reanalyze a possible origin of the resolved radio emission from various wind environment, and show that it presumably arises from jet(s) ejected from the hot component during its recent series of outbursts. ",
        "conclusions": "The major results and conclusions of this paper can be summarised as follows:  \\begin{description}  \\item ({i})  The radio emission from the symbiotic binary AG Dra seems to be variable, and probably related to the hot component activity.  \\item ({ii}) The cool component wind can, in principle, account for the intensity of the N1 component (O02), which position practically coincides with the {\\it Hipparcos} position of AG Dra.  \\item ({iii}) The two radio components resolved by MERLIN (O02) are practically aligned with the binary axis of AG Dra. The possible extended radio emission reported by Torbett \\& Campbell (1987) has similar orientation. The resolved radio emission presumably originates in jets ejected from the binary system.  \\item ({iv}) Assuming that the jet velocity is of order of the escape velocity of the hot component, the separation between the two radio sources indicates that the ejection took place $\\sim 3$ yr earlier, and it was associated with the recent series of outbursts.  \\end{description}  \\subsection*{ACKNOWLEDGEMENTS} I gratefully acknowledge very helpful comments on this project by M. Friedjung, A. Omont, and R. Viotti. This study was supported in part by the KBN Research Grant No 5P03D\\,019\\,20, and by the JUMELAGE program \"Astronomie France-Pologne\" of CNRS/PAN."
    },
    "0206/astro-ph0206014_arXiv.txt": {
        "abstract": "We revisit the issue of cosmological parameter estimation in light of current and upcoming high-precision measurements of the cosmic microwave background power spectrum. Physical quantities which determine the power spectrum are reviewed, and their connection to familiar cosmological parameters is explicated. We present a set of physical parameters, analytic functions of the usual cosmological parameters, upon which the microwave background power spectrum depends linearly (or with some other simple dependence) over a wide range of parameter values. With such a set of parameters, microwave background power spectra can be estimated with high accuracy and negligible computational effort, vastly increasing the efficiency of cosmological parameter error determination. The techniques presented here allow calculation of microwave background power spectra $10^5$ times faster than comparably accurate direct codes (after precomputing a handful of power spectra). We discuss various issues of parameter estimation, including parameter degeneracies, numerical precision, mapping between physical and cosmological parameters, and systematic errors, and illustrate these considerations with an idealized model of the MAP experiment. ",
        "introduction": "By January 2003, microwave maps of the full sky at 0.2$^\\circ$ resolution will be available to the world, the harvest of the remarkable MAP satellite currently taking data \\cite{MAP}.  The angular power spectrum of the temperature fluctuations in these maps will be determined to high precision on angular scales from the resolution limit up to a dipole variation; this corresponds to about 800 statistically independent power spectrum measurements. Recent measurements have given a taste of the data to come, although covering much smaller patches of the sky and with potentially more serious systematic errors \\cite{dib00,han00,hal02}.  The main science driving these spectacular technical feats is the determination of basic cosmological parameters describing our Universe, and resulting insights into fundamental physics; see \\cite{kk99,hd02} for recent reviews and \\cite{kos01} for a pedagogical introduction.  The expected series of acoustic peaks in the microwave background power spectrum encode enough information to make possible the determination of numerous cosmological parameters {\\it simultaneously} \\cite{jun96}. These parameters include the long-sought Hubble parameter $H_0$, the large-scale geometry of the Universe $\\Omega$, the mean density of baryons in the Universe $\\Omega_b$, and the value of the mysterious but now widely accepted cosmological constant $\\Lambda$, along with parameters describing the tiny primordial perturbations which grew into present structures, and the redshift at which the Universe reionized due to the formation of the first stars or other compact objects.  While most of these parameters and other information will be determined to high precision, one near-exact degeneracy and other approximate degeneracies exist between these parameters \\cite{bon94,bon99}.  The exciting prospects of providing definitive answers to some of cosmology's oldest questions raise a potentially difficult technical issue, namely finding constraints on a large-dimensional parameter space. Given a set of data, what range of points in parameter space give models with an acceptably good fit to the data? The answer requires evaluating a likelihood function at many points in parameter space; in particular, finding confidence regions in multidimensional parameter space requires looking around in the space. This is a straightforward process, in principle. But a parameter space with as many as ten dimensions requires evaluating a lot of models: a grid with a crude 10 values per parameter contains ten billion models. A direct calculation of this many models is prohibitive, even with very fast computers. While some of the parameters are independent of the others (i.e.  tensor modes), reducing the effective dimensionality, many of the parameters will be constrained quite tightly, requiring a finer sampling in that direction of parameter space. An additional problem with brute-force grid-based methods is a lack of flexibility: if additional parameters are required to correctly describe the Universe, vast amounts of recalculation must be done. Grid-based techniques have been used for analysis of cosmological parameters from microwave background (e.g., \\cite{gor94,teg00}) but clearly this method is not fast or flexible enough to deal with upcoming data sets adequately.  A more sophisticated approach is to perform a search of parameter space in the region of interest.  Relevant techniques are well known and have been applied to the microwave background \\cite{dod00,mel00,chr01}. Reliable estimates of the error region in cosmological parameter space can be obtained with random sets of around $10^5$ models, reducing the computational burden by a factor of 1000 or more compared to the cruder grid search methods. On fast parallel computers it is possible to compute the power spectrum for a given model in a computation time on the order of a second \\cite{spergel}, making Monte Carlo error determinations feasible. For improved efficiency, useful implementations do not recalculate the entire spectrum at each point in parameter space, but rather use approximate power spectra based on smaller numbers of calculated models \\cite{kno01}.  Here we expand and refine this idea by presenting a set of parameters, functions of the usual cosmological parameters, which reflect the underlying physical effects determining the microwave background power spectrum and thus result in particularly simple and intuitive parameter dependences. Previous crude implementations of this idea applied to low-resolution measurements of the power spectrum \\cite{dod95}; the current and upcoming power spectrum measurements requires a far more refined implementation. Our set of parameters can be used to construct computationally trivial but highly accurate approximate power spectra, and large Monte Carlo computations can then be performed with great efficiency.  Additional advantages of a physically-based parameter set are that the degeneracy structure of the parameter space can be seen much more clearly, and the Monte Carlo itself takes significantly fewer models to converge.  Earlier Fisher-matrix approximations of parameter errors are a rough implementation of this general idea: if the the power spectrum varies exactly linearly with each parameter in some parameter set, and the measurement errors are gaussian random distributed and uncorrelated, then the likelihood function as a function of the parameters can be computed {\\it exactly} from the partial derivatives of the model with respect to the parameters \\cite{jun96,zal97a,bon97}.  Even if the linear parameter dependence does not hold throughout the parameter space, it will almost always be valid in some small enough region, being the lowest-order term in a Taylor expansion. If this region is at least as large as the resulting error region, then the Fisher matrix provides a self-consistent approximation for error determination. The difficulty is in finding a suitable parameter set. Our aim is not necessarily to find a set of parameters which all have a perfectly linear effect on the power spectrum.  Rather, more generally, we desire a set of parameters for which the power spectrum can be approximated very accurately with a minimum of computational effort. Then we can dispense with approximations of the likelihood function, as in the Fisher matrix approach, and directly implement a Monte Carlo parameter space search with a minimum of computational effort.  This technique allows simple incorporation of prior probabilities determined from other sources of data, and if fast enough allows detailed exploration of potential systematic biases in parameter values arising from numerous experimental and analysis issues such as treatment of foregrounds, noise correlations, mapmaking techniques, and model accuracy.  We emphasize that with the accuracy of upcoming microwave background power spectrum measurements, the cosmological parameters will be determined precisely enough that their values will be significantly affected by systematic errors from assumptions made in the analysis pipeline and, potentially, numerical errors in evaluating theoretical power spectra. The measurements themselves may also be dominated by systematic errors. Extensive modelling of the impact of various systematic errors on cosmological parameter determination will be {\\it essential} before any parameter determination can be considered reliable. This is the primary motivation for increasing the efficiency of parameter error analysis. We discuss this point in more detail in the last Section of the paper.  The following Section reviews physical processes affecting the cosmic microwave background and defines our physical parameter set. The key parameter, which sets the angular scale of the acoustic oscillations, has been discussed previously in similar contexts \\cite{eis99,bon99,gri01} but has not been explicitly used as a parameter. The other parameters describing the background cosmology can then be chosen to model other specific physical effects.  Section \\ref{sec:power_spectra} displays how the power spectrum varies as each parameter changes with the others held fixed.  The mapping between our physical parameters and the conventional cosmological parameters immediately reveals the structure of degeneracies in the cosmological parameter space. Simple approximations for the effect of each physical parameter are shown to be highly accurate for all but the largest scales in the power spectrum.  We then determine the error region in parameter space for an idealized model of the MAP experiment in Sec.~\\ref{sec:errors}, comparing with previous calculations, and test the accuracy of our approximate power spectrum calculation. Finally, Sec.~\\ref{sec:discussion} discusses the potential speed of our method, likelihood estimation, mapping between the sets of parameters, accuracy, and systematic errors. The first appendix summarizes an analysis pipeline going from a measured power spectrum to cosmological parameter error estimates; the second appendix details several numerical difficulties with using the CMBFAST code for the calculations in this paper, along with suggested fixes.  Other recent attempts to speed up power spectrum computation through approximations include the dASH numerical package \\cite{kap02} and interpolation schemes \\cite{teg00}. Our method has the advantage of conceptual simplicity and ease of use combined with great speed and high accuracy. ",
        "conclusions": "\\label{sec:discussion}  Error regions similar to Fig.~\\ref{errors_cosmological_full} have been previously constructed for model microwave background experiments \\cite{bon99,lew02}. The remarkable point about our calculation is that the entire error region, constructed from power spectra for $3\\times 10^4$ points in parameter space, has been computed in a few seconds of time on a desktop computer. In fact, since our power spectra can be computed with a few arithmetic operations per multipole moment, the calculation time might not even be dominated by computing model power spectra, but rather by computing the likelihood function or by converting between the cosmological and physical parameters.  For actual non-diagonal covariance matrices describing real experiments, the computation of the likelihood will dominate the total computation time.  Gupta and Heavens have recently formulated a method for computing an approximate likelihood with great efficiency \\cite{gup01}, based on finding uncorrelated linear combinations of the power spectrum estimates \\cite{hea00}. Each parameter corresponds to a unique combination, so the likelihood is calculated with a handful of operations. Moreover, the method is optimized so the parameter estimation can be done with virtually no loss of accuracy. Using this technique, likelihood estimates for realistic covariance matrices will be roughly as efficient as our power spectrum estimates.  The conversion from the physical parameters to cosmological ones requires evaluation of several numerical integrals. The integrands will be smooth functions, and this step can likely be made nearly as efficient as the power spectrum evaluation, although we have not implemented an optimized routine for parameter conversion.  We make the following rough timing estimate: each multipole of the power spectrum can be approximated by a handful of floating point operations, and an approximate likelihood and the parameter conversion both will plausibly require similar computational work.  So we anticipate on the order of 10 floating point operations per multipole for each of 1000 multipoles.  On a 1 gigaflop machine, this results in evaluation of around $10^5$ models per second, so a Markov Chain with tens of thousands of models can be computed in under a second.  Of course, such a computation is trivially parallelizable, and can easily be sped up by a factor of tens on a medium-size parallel machine.  In contrast, other state-of-the-art parameter estimation techniques \\cite{kno01,kap02,spergel}, dominated by the power spectrum calculation, compute roughly one model per second, requiring several hours for a $10^4$ point computation.  This great increase in computational speed is highly useful.  All upcoming microwave background measurements and resulting parameter estimates will be dominated by systematic errors, both from measurement errors when observing the sky and processing errors in the data analysis pipeline.  The only way to discern the effect of these errors is through modelling them, requiring a determination of parameters numerous times. The microwave background has the potential to constrain parameters at the few percent level.  But a variation of this size in any given parameter will result in changes in $C_l$'s of a few percent, which is significantly smaller than the errors with which each $C_l$ will be measured.  This means that small systematic effects distributed over many Cl's can bias derived parameter values by amounts larger than the formal statistical errors on the parameter. Exhaustive simulation of a wide range of potential systematics will be required before the accuracy of highly precise parameter determinations can be believed, and our techniques make such investigations far faster and easier.  Along with increased speed, our approximation methods also promise high accuracy. A careful error evaluation is hindered by systematic errors in CMBFAST; our numerical approximations give better accuracy than the predominant numerical code for parameter ranges generally larger than the region which will be allowed by the MAP data. The physical set of variables presented here clarify the degeneracy structure of the power spectrum, clearly displaying degenerate and near-degenerate directions in the parameter space.  For example, with the physical parameters presented here, it is clear that any tensor mode contribution, which affects only low-$l$ multipoles, will have negligible impact on the error contours in Fig.~\\ref{errors_physical_full}.  Bond and Efstathiou used a principle component analysis to claim that uncertainties in the tensor mode contributions would be the dominant source of error for all of the cosmological parmeters, and further claimed that as a result, the Fisher matrix approximation will overestimate the errors in other apparently well-determined parameters such as our ${\\cal B}$ \\cite{bon99}. These results are clearly valid only for data which is significantly less constraining than the MAP data will be. Essentially, their approximate degeneracy trades off variations in $\\cal R$ and $\\cal B$, which change the heights of the first few peaks, with variations in $n_s$, which also impact the first few peaks, while holding $\\cal A$ fixed.  From our analysis, it is clear that this degeneracy only holds approximately, since $\\cal R$ and $\\cal B$ have a significant impact only on the first three peaks, while $n_s$ affects all of the peak heights.  Indeed, a careful examination of Fig.~3 in \\cite{efs02} shows that the power spectra are only approximately degenerate, with peak height differences which are significant for MAP. The addition of tensor modes allows a rough match between the two power spectra at low $l$ values, but for high $l$ values the discrepancy in peak heights is already evident at the third peak and will become larger for higher peaks, since the values of $n_s$ in the two models are so different.  Gravitational lensing makes a negligible difference in the linearity of the parameter dependences. We have not included polarization in this analysis, but the three other polarized CMB power spectra can be handled in the same way as the temperature case using standard techniques \\cite{kam97,zal97c}.  We have shown that a proper choice of physical parameters enables the microwave background power spectrum to be simply approximated with high accuracy over a significant region of parameter space. This region will be large enough for analyses of data from the MAP satellite, although larger regions could be stitched together using multiple reference models.  Our approximation scheme requires first computing numerical derivatives of the power spectrum multipoles with respect to the various parameters, but this must be done only once and then computing further spectra is extremely fast (in the neighborhood of $10^5$ models per second or more on common computers).  The error determination techniques in this paper require the numerical evaluation of only tens of power spectra rather than thousands or millions, so speed of the power spectrum code will not be of paramount concern. On the other hand, we emphasize that for {\\it any} high-precision parameter estimate, even small systematic errors in computing $C_l$'s can lead to biased parameter estimates comparable to the size of the statistical errors. These two considerations argue for a significant revision of CMBFAST (CMBSLOW, perhaps?)  or construction of other independent codes which focus on overall accuracy and stability of derivatives with respect to parameters rather than on computational speed. Such a code, combined with the estimation techniques in this paper and efficient likelihood evaluation methods \\cite{gup01} will provide a highly efficient and reliable way to constrain the space of fundamental cosmological parameters."
    },
    "0206/astro-ph0206045.txt": {
        "abstract": "{ Gamma-ray and microwave observations of the Cygnus region reveal an intense signal of 1.809 \\MeV\\ line emission, attributed to radioactive decay of \\al, that is closely correlated with 53 GHz free-free emission, originating from the ionised interstellar medium. We modelled both emissions using a multi-wavelength evolutionary synthesis code for massive star associations that we applied to the known massive star populations in Cygnus. For all OB associations and young open clusters in the field, we determined the population age, distance, and richness as well as the uncertainties in all these quantities from published photometric and spectroscopic data. We propagate the population uncertainties in model uncertainties by means of a Bayesian method. The young globular cluster Cyg OB2 turns out to be the dominant \\al\\ nucleosynthesis and ionisation source in Cygnus. Our model reproduces the ionising luminosity of the Cygnus region very well, yet it underestimates \\al\\ production by about a factor of 2. We attribute this underestimation to shortcomings of current nucleosynthesis models, and suggest the inclusion of stellar rotation as possible mechanism to enhance \\al\\ production. We also modelled \\fe\\ nucleosynthesis in the Cygnus region, yet the small number of recent supernova events suggests only little \\fe\\ production. Consequently, a detection of the 1.137 \\MeV\\ and 1.332 \\MeV\\ decay lines of \\fe\\ from Cygnus by the upcoming {\\em INTEGRAL} observatory is not expected. ",
        "introduction": "\\label{sec:intro}  OB associations and young open clusters constitute the most prolific nucleosynthesis sites in our Galaxy. The combined activity of stellar winds and core-collapse supernovae ejects significant amounts of freshly synthesised nuclei into the interstellar medium. Radioactive isotopes, such as \\al\\ or \\fe, that have been co-produced in such events may eventually be observed by gamma-ray instruments through their characteristic decay-line signatures. Indeed, galactic 1.809 \\MeV\\ gamma-ray line emission attributed to the radioactive decay of \\al\\ has been observed by numerous gamma-ray telescopes (see Prantzos \\& Diehl 1996\\nocite{prantzos96} for a review). In particular, the COMPTEL telescope provided the first image of the Galaxy in the light of this isotope, showing an asymmetric ridge of diffuse emission along the galactic plane with a prominent localised emission enhancement in the Cygnus region (Diehl et al.~1995\\nocite{diehl95}). The Cygnus emission has been interpreted as the result of \\al\\ ejection in Wolf-Rayet winds and during core collapse supernova explosions from a nearby (1-2 kpc) massive star population which probably is part of the local spiral arm structure (del Rio et al.~1996\\nocite{delRio96}).  Gamma-ray line emission is not the only tracer of this activity. \\cite{knoedl99a} demonstrated that the galactic 1.809 \\MeV\\ emission is closely correlated to galactic free-free emission as observed in the microwave domain. The free-free emission mainly results from the ionisation of the interstellar medium by the UV flux of O stars, hence it traces the massive star population. The observed correlation is indeed one of the strongest arguments in favour of prolific \\al\\ production by massive stars. In general, a wealth of distinct massive star populations of different ages, sizes, or metallicities contribute to the emission along a line of sight through the Galaxy, and the observed correlation allows only conclusions about the average properties of the contributing populations. In that way, the observations suggest that the equivalent O7V star \\al\\ yield, defined as the average amount of \\al\\ ejected per number of O7V star (measured by their equivalent ionising production), has a galaxywide constant value of $\\yal = (1.0 \\pm 0.3) \\times 10^{-4}$ \\Msol\\ (Kn\\\"odlseder 1999\\nocite{knoedl99}).  It is surprising, however, that the correlation between 1.809 \\MeV\\ and microwave free-free emission also holds for the Cygnus region. Both gamma-ray and microwave data show a localised emission enhancement towards Cygnus, similar in size and relative intensity, resulting in an equivalent O7V star \\al\\ yield of $(1.1 \\pm 0.3) \\times 10^{-4}$ \\Msol\\ (see Sect. \\ref{sec:results}). Within the uncertainties this yield is identical to the galactic value. In contrast to the Galaxy, however, only few massive star associations contribute to the observed emission in Cygnus, and it is not expected that they have the same properties as the Galaxy as a whole. In particular, the galactic metallicity gradient leads to an average galactic abundance that is supersolar, and indeed only a supersolar abundance is able to reconcile theoretical \\al\\ yields with the observed galactic \\al\\ mass (Kn\\\"odlseder 1999\\nocite{knoedl99}). In contrast, massive star populations in Cygnus show slightly subsolar abundances (e.g.~Daflon et al.~2001\\nocite{daflon01}) and since \\al\\ yields are believed to depend on metallicity (e.g.~Prantzos \\& Diehl 1996\\nocite{prantzos96}) the nucleosynthetic properties of the Cygnus region should deviate from those of the average Galaxy.  To understand the observations, we present in this paper a bottom-up model of the Cygnus region where we aim to explain the gamma-ray and microwave data from the underlying stellar populations. For this purpose we developed a multi-wavelength evolutionary synthesis model that we presented in paper I of this series (\\nocite{cervino00}Cervi\\~no et al.~2000). We put considerable effort into the characterisation of the massive star populations in Cygnus with particular emphasis on the involved uncertainties (distance and age uncertainty; coeval or continuous star formation). We incorporate these uncertainties into our model by means of a Bayesian method and determine confidence intervals for all quantities to assess the predictive power of our approach. Despite the resulting uncertainties, we will demonstrate that the gamma-ray observations provide important clues on nucleosynthesis physics in massive stars. In particular we will demonstrate the shortcomings of current theoretical nucleosynthesis models in explaining \\al\\ production and discuss possible modifications that may improve the models.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Evolutionary synthesis model and analysis method %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusions}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Nucleosynthesis} \\label{sec:nucleosynthesis}  Our modelling effort of the gamma-ray line emission from OB associations and young open clusters in the Cygnus region has revealed a possible shortcoming of actual nucleosynthesis models in explaining \\al\\ production. We find a 1.809 \\MeV\\ flux underestimation of about a factor of 2 that is difficult to explain by other means than a modification of current nucleosynthesis models for single, non-rotating stars. From preliminary calculations, it appears difficult for rotation alone to significantly enhance \\al\\ production, although this conclusion needs to be checked by more detailed calculations in the future.  Only little \\fe\\ production is predicted in Cygnus by our model, mainly related to the low number of recent supernova events in this region. A detection of the 1.137 \\MeV\\ and 1.332 \\MeV\\ lines from the radioactive decay of \\fe\\ by {\\em INTEGRAL} would therefore present a big surprise. In the case of such a detection, \\fe\\ production by hydrostatic helium burning in Wolf-Rayet stars -- which we have not included in our model due to the low yield predictions -- should then be seriously reconsidered (Arnould et al. 1997\\nocite{arnould97}).  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Other gamma-ray signatures} \\label{sec:signatures}  In addition to \\al\\ and \\fe, massive star associations may produce further radioactive isotopes during the supernova explosion of massive stars, such as \\co\\ and \\ti, yet their short lifetimes of $112$ days and $87$ years, respectively, make their observation impossible in absence of a very recent event. However, \\co\\ and \\ti\\ decay under positron emission (like \\al), and the annihilation of positrons with electrons of the interstellar medium on time scales of a few $10^5$ years may provide a reverberation of the short-lived, extinct, radioactivities. Hence, the observation of the 511 keV positron annihilation line may provide an independent measure of the supernova activity in massive star associations (and in particular in Cygnus), although the interpretation of the observations will be complicated by the annihilation physics, the positron transport, and the positron escape fraction from the expanding supernova remnants.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Super star clusters} \\label{sec:superstarclusters}  Gamma-ray as well as free-free emission in the Cygnus region seems to be dominated by a single, extremely massive association: Cyg OB2. Indeed, Cyg OB2 is merely a prototype of a young globular cluster than an OB association (Kn\\\"odlseder 2000\\nocite{knoedl00}). There are examples of further super star clusters in the Galaxy, such as NGC 3603 (Moffat et al. 1994\\nocite{moffat94}), the Arches and Quintuplet clusters near the galactic centre (Figer et al. 1999\\nocite{figer99}), or the W49A cluster (Conti \\& Blum 2002\\nocite{conti02}). Most of these super star clusters are partially or totally obscured in the visible by the absorbing effects of intervening and/or local interstellar dust, and it is unclear how many of them exist throughout the entire Galaxy. Making the simplifying assumption that super star clusters are distributed uniformly throughout the Galaxy within a galactocentric distance of 15 kpc, and taking that Cyg OB2 is apparently the only such object within 1.5 kpc, one may expect 100 super star clusters in our Galaxy. Assuming that they all produce $10^{-2}$ \\Msol\\ of \\al, similar to our Cyg OB2 model prediction, a total \\al\\ production of 1 \\Msol\\ is expected from these objects. Taking into account the \\al\\ yield underestimation of a factor of 2 brings this mass to 2 \\Msol, comparable to the observed galactic \\al\\ mass of $2-3$ \\Msol\\ (Diehl et al. 1995\\nocite{diehl95}; Kn\\\"odlseder 1999\\nocite{knoedl99}). Thus, a considerable fraction of \\al\\ maybe indeed produced by such super star clusters, and 1.809 \\MeV\\ gamma-ray observations with sufficient angular resolution and sensitivity may help to uncover and to study them throughout the Milky-Way.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{The Cygnus superbubble}  Our age determination of clusters in the Cygnus region sheds some doubt on the triggered star formation scenario that has been proposed to explain anomalous stellar proper motions in the area (Comer\\'on \\& Torra 1994\\nocite{comeron94}). In this scenario, the central association Cyg OB2 is supposed of having formed a shell blown by stellar winds and supernovae, the Cygnus superbubble, that subsequently gave birth to the surrounding associations Cyg OB1, OB3, OB7 and OB9 due to gravitational shell instability. However, our age estimate for Cyg OB2 is inferior to that of the surrounding associations, making triggered star formation in this area unlikely. In contrast, the anomalous stellar proper motions are equally well explained by supposing that they result from dynamically ejected runaway stars from Cyg OB2 (Comer\\'on et al. 1998\\nocite{comeron98}), in particular since the proposed expansion age of $\\sim4$ Myr is compatible with our age estimate for Cyg OB2.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Cygnus and the Galaxy} \\label{sec:galaxy}  Finally, we want to mention that not only the Cygnus region but also the Galaxy as a whole suffers probably from an underproduction of \\al\\ when current, non-rotating nucleosynthesis models are considered. By combining the same nucleosynthesis models that have been used in this work with the atmosphere models of \\cite{schaerer97}, and normalisation to the observed galactic Lyman continuum luminosity, \\cite{knoedl99} predicted a galactic \\al\\ mass of $1.6$ \\Msol, significantly below the observed value of $2-3$ \\Msol. Only the inclusion of the galactic metallicity gradient, which considerably enhances \\al\\ ejection through stellar winds in the Wolf-Rayet phase towards the galactic centre, provides a sufficient increase of the galactic \\al\\ production, bringing the model predictions in better agreement with the observations.  Yet, metallicity is of no help in the Cygnus region, where even slightly subsolar abundances are reported (Daflon et al. 2001\\nocite{daflon01}). On the other hand, if another process, such as rotation, is needed to explain the \\al\\ production in Cygnus, it should also be at work for the entire Galaxy, leading to a potential \\al\\ overproduction if combined with the metallicity enhancement. In fact, if \\al\\ production by massive stars would indeed follow the $(Z/Z_\\odot)^{2}$ dependence predicted by non-rotating stellar models, a break-down of the tight correlation between 1.809 \\MeV\\ and galactic free-free emission would be expected due to the inverse metallicity dependencies of both emission processes -- the 1.809 \\MeV\\ intensity should be enhanced in the metal-rich inner regions of the Galaxy due to enhanced \\al\\ production, while the free-free emission should be reduced due to lower electron temperatures (see Kn\\\"odlseder 1999\\nocite{knoedl99}).  The existing data show no indication for such an anti-correlation. However, it is likely that rotation reduces the metallicity dependence of Wolf-Rayet star yields. The fact that at low metallicity, there is less $^{25}$Mg, might be somewhat compensated by the fact that when the metallicity decreases, the mixing is more efficient, bringing more $^{25}$Mg from the radiative envelope into the core and more \\al\\ from the convective core into the radiative envelope. The mixing efficiency increases when the metallicity decreases because, when $Z$ decreases, the stars are more compact, the internal gradients of the angular velocity are steeper and the stars loose less angular momentum by mass loss through stellar winds (see Maeder \\& Meynet 2001\\nocite{maeder01}). Secondly, when rotation is accounted for, the mass loss rate plays a less important role in the WR formation process (see Fig.~\\ref{kipwral}). Thirdly, at low metallicity,  stars loose less angular momentum. As a consequence they can more easily reach the break-up limit at a given stage during their evolution. Very high mass loss rates ensue even if the metallicity is low. Finally, there are some indirect indications, that the proportion of fast rotators increases when the metallicity decreases (Maeder et al. 1999\\nocite{mgm}).  Since rotation seems to erode the metallicity dependence of Wolf-Rayet star yields, it provides an appealing mechanism to explain the absence of a 1.809 \\MeV\\ -- free-free emission anti-correlation. Remains to be seen if rotation can also be a solution to the \\al\\ yield puzzle -- both for the Cygnus region and the Galaxy as a whole.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Acknowledgements %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0206/astro-ph0206258_arXiv.txt": {
        "abstract": "We report results of $N$-body simulations of isolated star clusters, performed up to the point where the clusters are nearly completely dissolved. Our main focus is on the post-collapse evolution of these clusters. We find that after core collapse, isolated clusters evolve along nearly a single sequence of models whose properties are independent of the initial density profile and particle number. Due to the slower expansion of \\mbox{high-$N$} clusters, relaxation times become almost independent of the particle number after several core collapse times, at least for the particle range of our study. As a result, the dissolution times of isolated clusters exhibit a surprisingly weak dependence on $N$.  We find that most stars escape due to encounters between single stars inside the half-mass radius of the cluster. Encounters with binaries take place mostly in the cluster core and account for roughly 15\\% of all escapers. Encounters between single stars at intermediate radii are also responsible for the build up of a radial anisotropic velocity distribution in the halo. For clusters undergoing core oscillations, escape due to binary stars is efficient only when the cluster center is in a contracted phase. Our simulations show that it takes about $10^5$ $N$-body time units until the global anisotropy reaches its maximum value. The anisotropy increases with particle number and it seems conceivable that isolated star clusters become vulnerable to radial orbit instabilities for large enough $N$. However, no indication for the onset of such instabilities was seen in our runs. ",
        "introduction": "It is one of the great challenges of numerical astrophysics to understand the dynamical evolution of globular clusters and other collisional systems (as e.g.\\ galactic nuclei;  galaxy clusters). So far however, it is not possible to perform simulations of such rich systems by direct summation techniques and one has to rely on a statistical modeling. This is still the case, despite recent progress in the development of faster computers, like e.g.\\ the GRAPE computer series (Makino 2001, 2002).  Statistical modeling of globular clusters includes different techniques like Fokker-Planck, Monte Carlo and Gaseous model simulations. There has been substantial progress in recent years in improving the reality of such models and comparing them with $N$-body models or comparisons between different statistical methods show good agreement in many cases (Takahashi \\& Portegies-Zwart 2000, Giersz \\& Spurzem 2000). Nevertheless the validity of such codes for real globular clusters should be checked further, due to the complexity of the physical processes involved and the fact that the gap between the largest published $N$-body models (which are of order $3 \\cdot 10^4$), and the number of stars in real globular clusters (roughly $N = 10^6$) is still large.  Isolated star clusters are an ideal environment to test star cluster evolution. These systems are sufficiently simple, but nevertheless show many phenomena that are also seen in real globular clusters. Although simulations of isolated clusters are not directly applicable to real stellar systems, they are nevertheless highly relevant.  Borrowing an example from stellar evolution, such calculations play the role that polytropic models have played there historically.  While rather unrealistic physically, they did shed considerable light on the basic aspect of stellar structure, and in addition they served as useful comparison material when trying to interpret the results of the far more detailed numerical evolution models. Models for isolated clusters have the same dual role in stellar dynamics, they show some of the major physical effects occurring in long-term dynamical evolution, and provide templates against which more detailed models can be interpreted. A large amount of research has therefore been devoted to the study of isolated clusters.  It has been known for a long time that star clusters undergo a contraction of their center as a result of heat transfer from the cluster core to the halo (Antonov 1962, Lynden-Bell \\& Wood 1967). For isolated, single-mass clusters, this core collapse takes roughly 17 initial relaxation times until completion (Takahashi 1995), and the core collapse evolution is thought not to change very much in the presence of an external tidal field (Giersz \\& Heggie 1997). Lynden-Bell \\& Eggleton (1980) found that the core collapse proceeds from the outer to the inner parts and leaves behind a power-law debris of material with density exponent in the range $\\alpha = -2$ to -2.5. Although core collapse strongly influences the structure of the inner parts, one usually assumes that the overall properties of the system are determined more by the half-mass radius, and the core adjusts itself in post-collapse so as to balance the tendency of the halo to recollapse.  Core collapse will be reversed if a central energy source is present (H\\'enon 1975), and it is normally assumed that binary stars, which are either primordial or form during core collapse and harden as a result of encounters with field stars, provide this heat source (Aarseth 1971, Heggie 1975) in real star clusters. The number of binaries necessary to drive the cluster evolution is found to be quite small. Goodman (1984) estimated from his gaseous model calculations that their number drops with the number of cluster stars as $N_b \\propto N^{-0.3}$ and that in the mean only $N_b = 0.5$ binaries are necessary to drive the evolution of a cluster with $N = 10^6$ stars. This has been partly confirmed by Giersz \\& Heggie (1994b) in their $N$-body simulations of small-$N$ clusters. They found that only few binaries are present in the core, but their number seemed to increase with the particle number. They noted however that some of their binaries might be temporary ones which contribute nothing to the energy generation of the core.  Stars scattered out of the cluster center create an radially anisotropic velocity profile in the cluster halo (Larson 1970). Cohn (1984) found that a cluster starts to become anisotropic in its outer parts already early on in the pre-collapse phase, while Spitzer \\& Shapiro (1972) have argued that the anisotropy extends down into the core as a consequence of its collapse. This was later confirmed by Takahashi (1995) and Drukier et al.\\ (1999) for the late stages of core collapse. For post-collapse clusters, Takahashi (1996) found that they are isotropic in their center and the anisotropy increases monotonically towards the halo.  Encounters of stars in the cluster center also lead to the escape of stars. H\\'enon (1960) has shown that stars in isolated systems do not escape by distant encounters, but only due to close encounters with other stars. Despite their small numbers, it is generally believed that binaries play an important role in the production of escapers and contribute to the development of an anisotropic velocity profile in the cluster halo.  So far, most simulations of isolated star clusters were concerned only with the early stages up to core collapse or the immediate re-expansion phase. Nearly nothing is known about how clusters evolve when substantial mass-loss sets in. One reason for this neglect is that the escape of stars from isolated clusters happens very slowly and on a timescale which increases rapidly as the clusters expand. In addition, large-angle scatterings are neglected in Fokker-Planck methods and can only be studied by $N$-body or Monte Carlo simulations.  Hence, although highly desirable, it is difficult to verify the predictions of Fokker-Planck or Gaseous models by direct $N$-body simulations. We therefore present in this paper the first $N$-body simulations of medium sized isolated star clusters which follow the evolution until near complete disintegration. ",
        "conclusions": ""
    },
    "0206/astro-ph0206291_arXiv.txt": {
        "abstract": "In recent years, luminous X-ray outbursts with variability amplitudes as high as $\\approx 400$ have been serendipitously detected from a small number of active and inactive galaxies. These outbursts may result from the tidal disruptions of stars by supermassive black holes as well as accretion-disk instabilities. In order to place the first reliable constraints on the rate of such outbursts in the Universe and test the stellar tidal disruption hypothesis, we have performed a systematic and complete survey for them by cross-correlating \\rosat\\/ All-Sky Survey (RASS) and pointed PSPC data. We have detected five galaxies that were in outburst during the RASS, three of which show no signs of nuclear activity; these objects had been reported on individually in previous studies.  After making reasonable corrections for the complicated selection effects, we conclude that the rate of large-amplitude X-ray outbursts from inactive galaxies in the local Universe is $\\approx 9.1 \\times 10^{-6}$~galaxy$^{-1}$~yr$^{-1}$. This rate is consistent with the predicted rate of stellar tidal disruption events in such galaxies. When only the two active galaxies are considered, we find a rate for active galaxies of $\\approx 8.5 \\times 10^{-4}$~galaxy$^{-1}$~yr$^{-1}$. In order to place tighter constraints on these rates, additional outbursts must be detected. ",
        "introduction": "\\label{intro}  The detection of large-amplitude X-ray outbursts originating from inactive and active galactic nuclei (AGN) has generated considerable interest. These outbursts have variability amplitudes up to a factor of $\\approx 400$, decay over periods of months to years, often exceed outburst X-ray luminosities of $10^{43}$--$10^{44}$~erg~s$^{-1}$, and usually have extremely soft X-ray spectra.  Possible mechanisms considered to explain these events include (1) the tidal disruptions of stars by supermassive black holes, (2) accretion-disk instabilities, and (3) the X-ray afterglows of gamma-ray bursts (e.g., Komossa \\& Bade 1999 and references therein). Many of the outbursts detected thus far are best explained by the stellar tidal disruption scenario.  In addition, at least one of these outbursts appears to be consistent with the disruption of a brown dwarf or giant planet (Li, Narayan, \\& Menou 2002).  A typical inactive galaxy is expected to undergo a tidal disruption event as often as every $\\approx 10^{4}$--$10^{5}$~yr (e.g., Magorrian \\& Tremaine 1999). The resulting emission from such an event should peak in the extreme ultraviolet or soft X-ray band, and it should decline over a period of roughly several months (e.g., Gurzadyan \\& Ozernoy 1980; Rees 1990). Of course, the observable effects of these events remain fairly uncertain due to the complexity of the disruption and subsequent accretion processes; an alternate spectral distribution in which significant emission is radiated in the optical band has been investigated by Loeb \\& Ulmer (1997). Sembay \\& West (1993) predicted that at least several hundred and perhaps as many as several thousand tidal disruption events should have been detected during the \\rosat\\ All-Sky Survey if a substantial fraction of galaxies contain supermassive black holes (SMBHs) of masses $10^{7}$--$10^{8}$~\\msun; investigations of correlations between central black hole mass and bulge properties suggest that the latter is true (e.g., Ferrarese \\& Merritt 2000; Gebhardt et al. 2000). Detecting and understanding these events will aid in determining the importance of the stellar tidal disruption process in the fueling of SMBHs and may also lead to a better understanding of accretion-disk instabilities (e.g., Siemiginowska, Czerny, \\& Kostyunin 1996; Burderi, King, \\& Szuszkiewicz 1998). X-ray outbursts could also be related to the creation of double-peaked emission lines in AGN (e.g., Syer \\& Clarke 1992; Eracleous et al. 1995; Storchi-Bergmann et al. 1997) and nuclear outbursts seen at other wavelengths (e.g., Cappellari et al. 1999; Renzini et al. 1995).  The first large-amplitude X-ray outburst was detected in the galaxy E1615+061 using \\heao\\ and \\exosat\\ (Piro et al. 1988).\\footnote{We note, however, that a possible detection of a tidal disruption event was made in 1890 through the visual observation of the galaxy NGC~1068 (Packer 1891; de Vaucouleurs 1991).} All subsequent X-ray detections have been made using data from \\rosat.  \\rosat\\ was especially sensitive in the soft ($0.1$--$1.0$~keV) band where the emission from these outbursts is expected to be strong, and it covered $\\approx 20$\\% of the sky at least twice. Because constant monitoring of the X-ray sky or the comparison of at least two observations are the only ways to detect long-lived outbursts, \\rosat\\ provided an excellent means by which large-amplitude X-ray outbursts could be found and investigated.  Of those galaxies caught undergoing X-ray outbursts, only E$1615+061$, IC~3599 (Brandt, Pounds, \\& Fink 1995; Grupe et al. 1995a), and WPVS~007 (Grupe et al. 1995b) showed signs of nuclear activity prior to or after the outburst.  E$1615+061$ is a Seyfert 1, IC~3599 is a Seyfert 1.9 (Komossa \\& Bade 1999), and WPVS~007 is a Narrow-line Seyfert 1 (NLS1).  NGC~5905 (Bade, Komossa, \\& Dahlem 1996), RX~J$1242.6-1119$ (Komossa \\& Greiner 1999), RX~J$1420.4+5334$ (Greiner et al. 2000), and RX~J$1624.9+7554$ (Grupe, Thomas, \\& Leighly 1999) have optical spectra that show no signs of nuclear activity. Both NGC~5905 and RX~J$1624.9+7554$ are spiral galaxies; the former is classified as an HII-type.  RX~J$1242.6-1119$ is a pair of elliptical or early spiral galaxies likely to be interacting, and RX~J$1420.4+5334$ is also an elliptical or early spiral galaxy. It is likely that these four galaxies harbor otherwise dormant SMBHs that became active in the X-ray band only following transient fueling events.  The X-ray outbursts detected thus far were found either (1) serendipitously through X-ray and optical follow-up observations of galaxies or unidentified objects with soft X-ray spectra and other interesting properties, or (2) serendipitously from \\rosat\\ fields pointed at different targets.  In order to place reliable constraints on the number of such outbursts that occur in the Universe, more {\\it systematic\\/} and {\\it complete\\/} surveys must be performed.  Here we present the results of such a survey. By cross-correlating \\rosat\\ All-Sky Survey (RASS) and pointed observations, we have identified all \\rosat\\ sources at high Galactic latitudes that (1) were in outburst during the six-month RASS and that (2) had count rates or upper limits a minimum factor of $20$ lower in pointed observations taken before or after the RASS. We use the results of this survey to set the first reliable constraints on the frequency of large-amplitude X-ray outbursts in the Universe.  These constraints allow comparison with tidal disruption predictions such as those of Magorrian \\& Tremaine (1999), and they also allow assessment of the ability of \\chandra\\/, \\xmm\\/, and future missions to identify additional outbursts of this type.  We note that Komossa \\& Dahlem (2001) performed a similar survey for the nearby galaxies in the Ho, Filippenko, \\& Sargent (1995) sample, and did not detect any additional large-amplitude X-ray outbursts.  This work differs from ours in that X-ray outbursts were looked for only from previously known, nearby galaxies, whereas our survey looks for variability from all RASS sources.  Throughout this paper, values of $H_0= 75$~km~s$^{-1}$~Mpc$^{-1}$ and $q_0 = 0.5$ have been assumed. Galactic column densities have been taken from either Heiles \\& Cleary (1979) or Stark et al. (1992) as appropriate. ",
        "conclusions": "\\subsection{Constraints}  To place constraints on the rate of large-amplitude X-ray outbursts in the local Universe, we consider first the case in which all outbursts from galaxies and AGN have been included in the $\\log N$--$\\log S$ fit and for which the power-law index is fixed at $\\alpha = -3/2$. Although it is possible that the outbursts from active and inactive galaxies arise from different mechanisms, these events are still poorly understood, making such a calculation worthwhile.  The outburst rates derived from all fits can be found in Table~3.  For the fit described above, $\\approx 42$ outbursts should have occurred throughout the entire sky down to our characteristic completeness flux of $2 \\times 10^{-12}$~erg~cm$^{-2}$~s$^{-1}$. We assume a space density, $1.4 \\times 10^{-2}$~Mpc$^{-3}$, that is the sum of the inactive galaxy space density and the active galaxy space density. The inactive galaxy space density of $1.35 \\times 10^{-2}$~Mpc$^{-3}$ is the sum of a spiral galaxy space density of $1 \\times 10^{-2}$~Mpc$^{-3}$ (e.g., de Jong 1996) and an E+S0 space density of $3.5\\times 10^{-3}$~Mpc$^{-3}$ (e.g., Magorrian \\& Tremaine 1999). We assume an active galaxy space density of $5 \\times 10^{-4}$~Mpc$^{-3}$ (e.g., Peterson 1997).  The outburst rate for all galaxies and AGN is then $1.8 \\times 10^{-5}$~galaxy$^{-1}$~yr$^{-1}$. If intrinsic absorption is important, this rate will be higher; this effect is discussed in detail below for the case in which only outbursts from inactive galaxies are included in the $\\log N$--$\\log S$ fit.  When only the outbursts from inactive galaxies are considered, we calculate an outburst rate of $9.1 \\times 10^{-6}$~galaxy$^{-1}$~yr$^{-1}$. This rate corresponds to a timescale of $1.1\\times 10^{5}$~yrs between outbursts for a given inactive galaxy.  Because this rate lies just below the predicted maximum rates for tidal disruption events in inactive galaxies (see \\S 1), this result provides additional support for the hypothesis that large-amplitude X-ray outbursts from inactive galaxies are the result of stellar tidal disruptions by supermassive black holes.  In addition, this rate implies that down to 0.02~cts~s$^{-1}$, the typical count rate of a source is the \\rosat\\ All-Sky Survey Faint Source Catalog (RASSFSC)\\footnote{See http://www.xray.mpe.mpg.de/rosat/survey/rass-fsc/}, $\\approx 2000$ outbursts from inactive galaxies should have occurred during the RASS. This result is in agreement with the prediction of Sembay \\& West (1993; see \\S 1). We note that if intrinsic absorption prevents large-amplitude X-ray outbursts from being discovered, this rate will increase.  This effect was investigated in \\S 5 of Sembay \\& West (1993); they estimate that X-ray outbursts could be detected from only $\\approx 1/2$ of all spiral galaxies.  If we assume, from the respective number densities of spiral and E+S0 galaxies, that $\\approx 40$\\% of the inactive galaxies sampled by our survey are such roughly edge-on spirals, an outburst would be detectable from only $f \\approx 60$\\% of the galaxies in our sample. Consequently, the outburst rate would rise to $1.5 \\times 10^{-5}$~galaxy$^{-1}$~yr$^{-1}$.  Given the uncertainty on $f$, however, we perform all calculations under the assumption that $f$ is equal to unity.  If we consider only the two active galaxies in our sample for the case in which $\\alpha = -3/2$, we obtain an active galaxy outburst rate of $8.5\\times10^{-4}$~galaxy$^{-1}$~yr$^{-1}$. This rate is substantially higher than that for inactive galaxies. Because of the theoretical uncertainties associated with the tidal disruption of a star in an active galaxy, predictions of the rate of such events have not yet been made.  Consequently, it is difficult to determine whether the derived outburst rate for active galaxies is consistent with the stellar tidal disruption scenario.  The dynamics of a stellar tidal disruption event are expected to differ for a galaxy in which an accretion disk is present. It is thought that interactions between a star and the accretion disk of an active galaxy may allow the star to lose momentum and energy and reach a radius at which tidal disruption could occur (e.g., Syer, Clarke, \\& Rees 1991; Armitage, Zurek, \\& Davies 1996); the effect that this interaction would have on the rate of disruption is not yet clear.  The high apparent rate of outbursts from AGN, however, suggests that at least some of these outbursts may be due to another mechanism, such as accretion-disk instabilities (e.g., Siemiginowska et al. 1996; Burderi et al. 1998), which could lead to variations in the soft excess.   \\subsection{Predictions for Future Work}  In order to place tighter constraints on the outburst rates of both active and inactive galaxies, additional outbursts need to be detected.  It would also be useful if the outbursts could be followed-up in order to (1) determine their decay rates, (2) measure their X-ray spectra, and (3) look for associated spectroscopic signatures in other bands.  In the future, cross-correlations of \\rosat, \\xmm, and \\chandra\\/ data may provide a means by which to identify further outbursts.  Assuming that \\xmm\\/ operates for $\\approx 10$~yrs at 70\\% efficiency, taking observations with a mean length of 40~ks, it will produce $\\approx 5500$ observations.  If $\\approx 1000$ of these observations overlap and $\\approx 30$\\% are in the Galactic plane, \\xmm\\/ should produce $\\approx 3200$ distinct extragalactic observations.  Of these observations, $\\approx 35$\\% should have column densities $\\geq 3.8 \\times 10^{20}$~cm$^{-2}$, the highest column density through which an outburst in our survey was detected.  Consequently, $\\approx 2100$ observations will be useful for identifying large-amplitude X-ray outbursts.  The EPIC PN camera, which is the \\xmm\\/ instrument most sensitive to soft X-rays, has a field of view of $\\approx 718$~arcmin$^{2}$.  Approximately 1/2 of this area, $\\approx 359$~arcmin$^{2}$, is of sufficient sensitivity to detect these outbursting objects in quiescence for a 40~ks exposure. Consequently, \\xmm\\/ will cover $\\approx 0.5$ \\% of the extragalactic, low Galactic column density sky.  Cross correlation of these data with the RASS catalogs should allow the identification of all RASS outbursts with outburst count rates $\\geq 0.02$~cts~s$^{-1}$, the typical count rate of a source in the \\rosat\\/ All-Sky Survey Faint Source Catalogue (RASSFSC).\\footnote{See http://www.xray.mpe.mpg.de/rosat/survey/rass-fsc/} A count rate of 0.02~cts~s$^{-1}$ corresponds to an outburst flux of $\\approx 9\\times 10^{-14}$~erg~cm$^{-2}$~s$^{-1}$.  Consequently, cross-correlation between RASS and \\xmm\\/ observations should result in the detection of $\\approx 22$ RASS outbursts. We have carried out a similar calculation for \\chandra\\/ and find that, if \\chandra\\/ were highly sensitive in the 0.2--0.3~keV band, it could identify $\\approx 13$ RASS outbursts. The back-illuminated and front-illuminated ACIS CCD chips, however, are sensitive down to $\\approx 0.3$~keV and $\\approx 0.5$~keV, respectively.  Because the flux of a typical large-amplitude X-ray outburst drops by a factor of $\\approx 10$ when the flux from 0.2--0.3~keV is excluded, it will be significantly more difficult to prove variability by a factor of 20 without sensitivity in this energy range.  Consequently, only a handful of RASS outbursts could be expected to be detected through cross-correlations of \\rosat\\/ and \\chandra\\/ data.  Any correlation between all-sky survey and pointed observations, or pointed observations and pointed observations, will be limited by the complicated selection effects discussed in this paper, and perhaps others as well. As such, all-sky monitors, or a new sensitive all-sky survey, would provide the best means by which to detect and study additional outbursts in a uniform manner.  Missions such as \\lobster\\ (e.g., Parmar 2001), \\maxi\\ (e.g., Mihara et al. 2000), \\rosita\\footnote{See http://wave.xray.mpe.mpg.de/rosita}, and \\swift\\ (e.g., Gehrels 2000), however, probably do not have the combined spatial resolution and sensitivity to soft X-rays needed to identify additional outbursts of this type (see Grupe 2002). As such, a new sensitive soft X-ray all-sky survey or all-sky monitor is needed. Based on our fixed-$\\alpha$ $\\log N$--$\\log S$ fit for all outbursts from galaxies and AGN, we find that down to a count rate of $0.02$~cts~s$^{-1}$, the typical count rate of a source in the RASSFSC, $\\approx 4400$ large-amplitude X-ray outbursts should have occurred during the RASS.  Although the Galactic column density toward $\\approx 65$\\% of these sources is higher than the maximum Galactic column through which one of the outbursts in our sample was detected, $\\approx 1540$ outbursts should have occurred in regions of sky where the Galactic column density is sufficiently small. In order to detect and study these events, a second all-sky survey sensitive in the 0.1--2.4~keV range is needed.  Such a survey should be deep enough to detect sources with fluxes of $\\approx 5 \\times 10^{-15}$~erg~cm$^{-2}$~s$^{-1}$, a factor of 20 below the flux of an outburst with a 0.1--2.4 keV RASS count rate of $0.02$~cts~s$^{-1}$. Cross-correlation of the RASS with a new all-sky survey would remove many of the complicated selection effects considered above and would provide an excellent means to increase greatly the number of known outbursts of this type, allowing tighter constraints to be placed on the rate of these outbursts in the Universe. A sensitive soft X-ray all-sky monitor would provide an excellent means to detect newly outbursting objects and study them as their fluxes decline. Observations of possible tidal disruption events and accretion-disk instabilities in other wavelength bands will also be of use in understanding better these events."
    },
    "0206/astro-ph0206402_arXiv.txt": {
        "abstract": "This paper reports statistically significant correlations between various burst parameters, observed in a sample of 156 GRBs belonging to BATSE 4B catalog with T90 less than 2 s. The number of subpulses in a burst is strongly correlated not only with the object's duration but also with its fluence and hardness ratio, suggesting that when the central engine is more powerful, ejecting matter with typically higher values of Lorentz factor, the bulk energy is dissipated on longer time scales in the form of larger number of gamma pulses. We estimate hard-to-soft lag in bursts by taking the difference between centroids corresponding to time profiles at energies $> 100$ keV and $<100$ keV. The number of short GRBs that show soft-to-hard spectral evolution is slightly over one quarter  of the total, in the sample considered here.  Bursts that exhibit hard-to-soft spectral change appear to form a distinct class, with strength as well as hardness of individual subpeaks tending to decrease with peak position. Opposite is true for objects with softer photons arriving earlier than the harder ones, implying some kind of a rejuvenation of the central engine (may be due to enhanced accretion of matter towards the end). The two classes also show other diverging trends. For instance, objects belonging to the larger of the two classes display strong correlations between spectral lag and the fluence, the hardness ratio as well as the number of pulse, respectively. While no such correlations are seen in bursts that evolve from soft to hard. However, the magnitude of lag is strongly correlated with burst duration in both the classes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206416_arXiv.txt": {
        "abstract": "Available data on the magnesium - velocity dispersion (Mg-$\\sigma$) relation for $\\sim$2000 early-type galaxies is collected and compared. As noted previously, the Mg residuals from a fitted line are roughly Gaussian near the mean but have an asymmetric blue tail, probably from subpopulations of relatively young stars. We define statistics for scatter and asymmetry of scatter in the Mg dimension and find impressive uniformity among data sets. We construct models of galaxy formation built to be as unbiased as possible toward the question of the importance of mergers in the formation of early type galaxies. The observational constraints (Mg-$\\sigma$ width, asymmetry, and mean Mg strength, plus mean age and width of abundance distribution) are severe enough to eliminate almost all of models. Among the casualties are: models with merger rates proportional to $(1+z)^{n}$ with $n>0$, models that assume early formation followed by recent drizzling of new stars, merger-only models where the number of mergers exceeds $\\approx$80, merger-only models with less than $\\approx$20 mergers, and models with a cold dark matter power spectrum even with biasing included. The most successful models were those with merger probability constant with time, with the number of mergers needed to form the galaxy around 40. These models are characterized by mean light-weighted ages of 7-10 Gyr (consistent with spectroscopic studies), a narrow abundance distribution, and a lookback time behavior nearly indistinguishable from passive evolution of old stellar populations. ",
        "introduction": "The process of formation of early-type galaxies is still largely unknown. Under study since \\citet{Larson(1975)}, the choice between formation of elliptical galaxies by a single brief collapse or via merging is still disputed in the astronomical literature.  Some properties of early-type galaxies are very uniform, suggesting a homogeneous and probably ancient origin. Elliptical and S0 galaxies have smooth light profiles and very low gas and dust content, like overgrown globular clusters. As a class, they display scaling relations such as a tight Fundamental Plane (FP: \\citet{djor87, dress87}) among the variables luminosity, size, and velocity dispersion, a narrow Mg-$\\sigma$ relation between integrated starlight absorption line strength and velocity dispersion \\citet{bend93, Zieg97}, and an orderly color-magnitude relation.  Many authors note that elliptical galaxies nearby and at intermediate redshift have properties consistent with old, passively evolving systems \\citep{Stand98,vandok96,kel97,ellis97,koda97,koda98,bend98, vandok98}.  Confusingly, there is also a great deal of evidence that favors a messier, ongoing formation process for early-type galaxies. Some galaxy mergers are caught in the act \\citep{toom77} with the theoretical expectation that the remnant will soon relax to resemble an elliptical, with an exponential light profile. The tails, ripples, shells, and other morphological aftereffects of merging are seen in many nearby elliptical galaxies \\citep{sch1, sch2, goud01, barb02, mark02}.  Measurements of the integrated stellar absorption features compared with stellar population models indicate light-weighted mean ages for the near-nuclear regions of elliptical galaxies that range from ancient through a median of $\\sim$7 Gyr to a few that appear less than 1 Gyr old \\citep{gonza93,wor97,terle02}. Young stellar populations are much brighter than old populations, so that a relatively modest burst of young stars may be able to skew the mean age of an essentially old galaxy to appear much younger than a mass-weighted mean age, but the presence of any youthful subpopulation contradicts the uniform-and-old hypothesis.  With this paper we attempt to bring clarity to the issue.  To simplify, we consider only the relation between the central velocity dispersion $\\sigma$ and the strength of the integrated stellar Mg and MgH features around 5100\\AA. The advantages of this relation are its distance independence and its small scatter \\citep{jorg96,bend98}. One disadvantage is that age and metallicity both have a similar effect on the Mg strength \\citep{wor94,forb01,coll99}: an older galaxy has a stronger Mg feature, but Mg strength can also be increased by Mg abundance. It is clear that the origin of the relation itself is primarily one of abundance: larger galaxies have more heavy elements. Intrinsic scatter exits in the $Mg_2$-$\\sigma_0$ relation ($\\sigma_0$ refers to an aperture-corrected central velocity dispersion). This scatter is not correlated with the appearance of the galaxies, with the degree of velocity anisotropy, with their environment \\citep{dress87,burs90}, or with deviation of the objects within (or perpendiculiar) to the FP. \\citet{bend93} intrepreted this as a combination of age and/or metallicity differencies among ellipticals with similar luminosity, and \\citet{coll99} find that both mean age and metallicity vary, probably in a correlated manner with younger-appearing galaxies tending to be somewhat more metal-rich. This correlation is also seen with Balmer-metal indicators \\citep{wor95,terle02}  The residuals of $Mg_2$ have a Gaussian core, but the galaxies with weaker $Mg_2$ at given $\\sigma_0$ form a tail. That is, the residual distribution is skewed with an excess of Mg-weak galaxies \\citep{bend93}. This is suggestive of late star formation: A youthful galaxy will have weak Mg strength from the presence of bright, hot stars with weak Mg strength and will appear as a low outlier.  We test two opposing hypotheses. In both cases, the Mg-$\\sigma$ relation is a mass-abundance correlation. In the ``ancient formation'' hypothesis, elliptical galaxies formed at the epoch of galaxy formation at high redshift (we assume $z_f$=5). The Mg-$\\sigma$ relation was put in place at that time. Furthermore, the symmetric portion of the scatter was put in place at formation. (We could be even more extreme and suggest that the asymmetric part of scatter was also imprinted at formation by an asymmetric abundance distribution, but the implication of this is trivial: except for the effect of passive stellar evolution the Mg-$\\sigma$ relation would remain unchanged back to $z_f$.) In the ``ancient formation'' hypothesis, mergers, as inconsequential as possible, cause occasional bursts of star formation that drive a few galaxies to weak Mg strength and cause the asymmetry of the observed relation. In the ``maximum merger'' hypothesis, the underlying Mg-$\\sigma$ relation is very tight, and all scatter, symmetric or not, is caused by merging events (convolved with observational error).  To compare these two views, we first collect and analyze published Mg-$\\sigma$ data sets of E and S0 galaxies from field, group, and cluster environments. We use these data to find robust statistics to help us evaluate the expected asymmetrical scatter. In $\\S$3 we describe special-purpose models that predict Mg-$\\sigma$ relations given input parameters such as the number of mergers and gas fraction of each merger. Results are discussed in $\\S$4, and a summary of conclusions is given in $\\S$5. ",
        "conclusions": "\\subsection{Basic Model Behavior}  As an introduction to the model output, we focus on examples of the effect of our four main parameters.  Figure~\\ref{fig3-a} shows the effect of varying the number of equal-mass mergers that eventually form the galaxy. With a lot of mergers (80 mergers, for example), the asymmetry becomes low since all the galaxies have almost the same history. That is, they all suffer a peppering of small amounts of star formation over the age of the universe.  But with very few mergers (5 mergers, for example) the dispersion is larger and the asymmetry is much more pronounced. Each galaxy experiences a distinct timing fingerprint for mergers and thus are more individualistic, some having had recent star formation, others having had none.  \\placefigure{fig3-a} \\placefigure{fig3-b} \\placefigure{fig3-c} \\placefigure{fig3-d} \\placefigure{fig4}  The shape of the probability envelope has a strong impact as well. The frequency of mergers is often parameterized in the literature as being proportional to $(1+z)^n$, where $n$ is a free parameter. Figure~\\ref{fig4} translates some of these choices from redshift to time.  One can see that $n = -1.2$ roughly corresponds to a uniform probability with time (note that when we say ``uniform probability'' throughout this paper, we mean a perfectly flat [with time] curve, avoiding altogether the $(1+z)^n$ formalism.); $n = -0.7$ corresponds to a present-day merger frequency about half of that at $z=z_f$, and $n > 0$ corresponds to merger rates strongly skewed to high redshift. These numbers change somewhat with cosmology.  As fig~\\ref{fig3-b} shows, it is plain that with $n = 0.5$ the obtained distribution is very thin. Since mergers are concentrated in the past, few are lucky enough to have had recent star formation. The relation thickens with decreasing $n$.  The parameter $F_{gas}$ is the fraction of gas used in each merger event. In Figure~\\ref{fig3-c} we see that more gas leads to more star formation and hence more young galaxies. Less gas leads to weak starbursts and a narrow Mg-$\\sigma$ relation.  The fraction of galaxy formed at $z=z_f$, $F_{ini}$, has less of a dramatic effect (see fig~\\ref{fig3-d}) except for the trivial case of $F_{ini}\\approx 1$. There is little difference between forming half of the galaxy at $z_f$ and forming the galaxy entirely by mergers, rather to the initial surprise of the authors. Apparently, it is the last few Gyr of history that are the most important.  \\subsection{Maximum Merger Models}  We now search quantitatively for solutions under the maximum merger hypothesis.  For fair statistical comparison, the artificial data were further randomized by an artificial Gaussian scatter $\\sigma_a = 0.01$ mag, representing the observational error only. That is, $\\sigma_a < \\sigma_g$, the width of the Gaussian core of the residual distribution. This assumption is modified for the ``ancient formation'' hypothesis described below.  Because we compare to the mean of all of the observational data, we analysed data from a similar number of artificial galaxies: roughly 2000 Monte Carlo galaxies. The scatter from limited sample size of course decreases with increased number of samples: for 250 galaxies the errors were .0009 and .2 for respectively the ADev and the ADR (computed from multiple realizations of the same simulation). With 2000 galaxies we obtained .0002 for the ADev and .04 for the ADR. These errors are not too different from those seen in the observational data sets, so it is all but certain that sampling error causes most of the statistical deviations between data sets. This leads to important guidelines for future work at higher redshift: large sample size is crucial if the scatter is to be adequately characterized, and errors may not creep too much higher than their nearby brethren.  \\placefigure{fig5-a} \\placefigure{fig5-b} \\placefigure{fig5-c}  Figures \\ref{fig5-a} through \\ref{fig5-c} show our main diagnostic statistics, ADev and ADR, for models that cover the grid of input parameters. Data for a uniform probability are shown in Figure~\\ref{fig5-a}, for a $(1+z)^0$ envelope in Figure~\\ref{fig5-b}, and for a $(1+z)^{-0.5}$ envelope in Figure~\\ref{fig5-c}.  Different $F_{ini}$ and $F_{gas}$ possibilities are plotted as various symbol types as a function of the number of mergers. The trend with increasing number of mergers is clear: a tighter, more symmetrical relation is found. If this seems counterintuitive because with many mergers one should expect a more volatile and bursty collection of galaxies, consider instead that each galaxy has a more similar history if the number of mergers grows large. See also Figure \\ref{fig3-a}. We plot a good sampling of the possibilities, representing F$_{ini} = 0.0$ or $0.5$ and F$_{gas} = .1, .2, .5$ with different symbol types. To compare directly with the observational data, we plot the mean observational values found in $\\S2$ with $\\pm 1.0\\,\\sigma$ and $\\pm 2.5\\,\\sigma$ errors.  Other trends are also apparent. At a given F$_{ini}$, the ADev increases with the gas fraction since with more gas available more stars will form, making the mean stellar population more volatile. Similarly, at fixed gas fraction, the ADev decreases with higher $F_{ini}$. Indeed, if $F_{ini}$ is set to one, the relation becomes razor thin since the whole galaxy formed at $z=z_f$. Except for values approaching one, the authors were surprised at the ineffectiveness of $F_{ini}$ to modify the Mg-$\\sigma$ relation by much; it is the weakest of the four parameters.  As for the influence of when in time the mergers are likely to happen, we compare the three figures. It is rather difficult to find many points that land simultaneously in the 2.5-$\\sigma$ regions of both ADev and ADR statistics, but to the extent that some do, most of them are in the ``uniform probability'' figure, around 40 mergers. Any probability exponent $n$ greater than zero yields Mg-$\\sigma$ relations that are even thinner than the $n=0$ case, and we are forced to conclude that elliptical galaxy formation that was concentrated to the first half of a Hubble time is ruled out under the maximum merger hypothesis.  \\subsection{Ancient Formation Models}  We now relax our assumption that the Mg-$\\sigma$ relation is inherently thin. We suppose instead that there is a intrinsic, symmetric scatter $\\sigma_g$ equal to that which is observed today, but imprinted at formation. The mean width derived in $\\S$2 is $\\sigma_g = 0.015$ mag. This is rather narrow and not too different from the observational scatter, 0.010 mag. In fact, consideration of these numbers implies that the true galaxy-intrinsic Gaussian-random scatter is around 0.01 mag. It will be interesting to see if future high-precision data upholds this conclusion.  Under the ancient formation hypothesis, only the asymmetric tail of the distribution is produced by late star formation. Do these models fit better? We plot the three different probability cases exactly as in the previous section: uniform probability (fig~\\ref{fig6-a}), $(1+z)^0$ envelope (fig~\\ref{fig6-b}) and $(1+z)^{-0.5}$ envelope (fig~\\ref{fig6-c}). We derived the statistics using 2000 galaxies, and show the possibilities allowed by $F_{gas} = 0.1,0.2,0.5$ and $F_{ini} = 0.0$ or $0.5$.  \\placefigure{fig6-a} \\placefigure{fig6-b} \\placefigure{fig6-c}  We find that most of the remarks from last section are still valid. The cases with 5 or 10 mergers do not fit, and the most numerous solutions still occur for uniform probability and $N_{\\rm merger}=40$.  The main difference is that the additional artificial scatter calms the asymmetry somewhat, and this allows a few more successful models. We even find a couple of models in the $n=0$ panel that land within the error bars for both statistics, although mostly we find the width of the Mg-$\\sigma$ relation too narrow.  The astounding similarity between maximum-merger and minimum-merger cases is driven by the high quality and consistency of the observational data. With the low and high bounds of allowed artificial scatter set at 0.010 and 0.015 mag, tight reigns are held on our freedom to interpret the results.  \\subsection{Constraints from Mg$_{300}$}  To this point in the paper we have examined only 'primary' statistics ADev and ADR for the width and asymmetry of the Mg-$\\sigma$ residuals. We can explore the situation a little more by also considering the absolute Mg strength of the Mg-$\\sigma$ relation, parameterized by Mg$_{300}$, the median Mg strength at $\\sigma =$ 300 km/s. Including Mg$_{300}$ along with ADR and ADev, we define a goodness parameter to evaluate the relative merit of the various models: \\begin{displaymath} \\sigma^2_G = \\frac{(ADev_i-\\overline{ADev_{obs}})^2}{\\sigma^2_{ADev}} + \\frac{(ADR_i-\\overline{ADR_{obs}})^2}{\\sigma^2_{ADR}} + \\frac{(Mg_{300_i}-\\overline{Mg_{300_{obs}}})^2}{\\sigma^2_{Mg_{300}}} \\end{displaymath}  We took $\\overline{ADev_{obs}} = .0206$ , $\\overline{ADR_{obs}} = 1.265 $ and $\\overline{Mg_{300_{obs}}} = 0.329$, for the different scatters we had: $\\sigma_{ADev} = 1.4\\,10^{-3} $, $\\sigma_{ADR}= 5.8\\,10^{-2} $ and $\\sigma_{Mg_{300}} = 0.04$.  This last value is less well-motivated than the others since we expect considerable uncertainty in model Mg zeropoint so that cannot adopt the observed uncertainty for $\\sigma_{Mg_{300}}$. On the other hand, the overall Mg level is an important diagnostic and too-weak overall Mg is clearly to be discouraged, so our value of 0.04 is an approximate compromise.  We compute this $\\sigma_G$ for the usual 3 cases : uniform merger probability envelope, (1+z)$^0$ and (1+z)$^{-0.5}$ probability, and for both F$_{ini} = 0$ or $0.5$. For each case, we make plots from 5 to 100 mergers, and from F$_{gas} = 0$ to F$_{gas} = 0.9$, for 40 values using a linear distribution for the fraction of gas and a logarithmic distribution for the number of mergers. We plot isovalues in 1/$\\sigma_G$ so that the best models show as maxima rather than minima. Fig~\\ref{fig7} shows the results for uniform, (1+z)$^0$ and (1+z)$^{-0.5}$ probability envelopes.  \\placefigure{fig7}  In this plot there are elongated regions of high $\\sigma_G$ values. They tend to slant, indicating significant degeneracy between gas fraction and number of mergers in determining a successful model. The amplitude of the band decreases with the $n$ of the envelope probability. We see that the valid possibilities shift to slightly higher values of gas fraction when F$_{ini} = 0.5$. The uniform envelope probability appears to be the best fit.  \\subsection{Constraints from Mean Age and Abundance Spread}  We now check to see if the possibly successful models indicated by ADev, ADR, and Mg$_{300}$ in Figure \\ref{fig7} are compatible with the abundance distribution and the V-band light-weighted age observed in nearby galaxies. The abundance distribution is observed to be narrow in all galaxies the size of M32 and larger \\citep{wor96}, with a FWHM of less than 0.4 dex \\citep{grill96} in M32 itself, and about the same in the Milky Way (e.g. \\citet{rana91}). We can check the abundance distributions of the simulations as well. The chemical evolution of these models is primitive, but if the abundance distributions are wider than about 0.5 dex FWHM then we must regard them with suspicion. Also, if the light-weighted mean age (the closest our present models can come to approximating a spectroscopic Balmer-metal feature mean age) falls below about 7 Gyr (the approximate median of such studies, e.g. \\citet{terle02, gonza93}) we likewise regard this model with much suspicion. Older ages are allowed because the age zeropoint of the Balmer-metal feature technique is still quite uncertain.  What we find is emphatic. For $n=0$ or $n=-0.5$ models, no model passes the ADR/ADev test very well, but they also tend to have an abundance scatter greater than 0.5 dex. For uniform probability, the mean galaxy age begins to becomes too young for a combination of high gas fraction and large numbers of mergers, but this never became a critical problem because the ADR/ADev test excludes these models as well. The best surviving model was with uniform probability, $F_{ini}=0.5$, $F_{gas}=0.35$, and 50 mergers, with an abundance histogram at the 0.5 dex width limit, mean age of 8.3 Gyr, and ADR and ADev within the error bars. At large number of mergers, the $F_{ini}=0.5$ models tend to have a bimodal abundance distribution caused by the initial metallicity of the galaxy, to which is added additional material from the gaseous portion of the merger, plus a low-metallicity peak from the stellar component of the relatively low-mass merger fragments. In fact, all of the $F_{ini}=0.5$ high points in Fig. \\ref{fig7} are at the borderline of having a too-wide abundance distribution, and we should penalize these models a couple of contour levels for this near-transgression. The $F_{ini}=0$ models do not suffer from this: they all have suitably narrow abundance distributions. We claim roughly equal success for both versions of the uniform-probability models, but resounding failure for models in which galaxy mergers happen preferentially in the past.  \\subsection{Model Variants}  Being left with only one fully successful model after trying so many hundreds left the authors scratching their heads. We decided to explore other options. Our first option was to abandon the constant merger mass assumption by adopting a cold dark matter (CDM) power spectrum for galaxy merger fragments.  We use the formulae from \\citet{white91} and \\citet{kauf92} which are based on \\citet{press74} formalism. CDM clusters hierarchically with preferentially small masses in the early universe, building toward larger structures later on. We built a lookup table with 40 steps in redshift and 200 steps in mass in order to quickly invert the probability function for halos of a given mass and redshift. Within each step a small randomization assigned the final redshift and mass to the merger fragment. To calibrate the zero point between the circular velocity (V$_c$, which the White \\& Frenk formulae give) and mass, we normalized to the Milky Way.  We took V$_c =$ 220 km/s, mass equal to 10$^{11}$ M$_{\\odot}$, and White \\& Frenk formulae to obtain M$_{gal =}$ 10$^{11}\\, (\\frac{V_c}{220\\, km/s})^3$.  The only parameters in the CDM merger models were gas fraction and the CDM ``bias'' parameter $b$. The ADR statistics from the CDM models were very high unless the gas fraction was set below 10\\%, at which point the ADev becomes too narrow. The $V$-weighted age and abundance scatter for the $F_{gas}=0.1$ model is acceptable: 11.4 Gyr and 0.4 dex, respectively.  Increasing the bias parameter helps a little, but no CDM models were found that matched all statistics simultaneously. The CDM simulations were also interesting because they predict variations in Mg-$\\sigma$ as a function of galaxy size. For small galaxies compared to large, CDM predicts similar ADev, younger mean ages, and much higher asymmetry. Spectroscopic observations \\citep{terle02} indicate a younger mean age for smaller galaxies. Our collection of Mg-$\\sigma$ observations indicate higher ADev but ADR roughly the same.  A variant on the CDM models was to throw out the CDM power spectrum and adopt a constant power spectrum ($dN \\propto M^{-1}$) and return to a merger probability uniform with time. These models have suitable ADR and ADev if the gas fraction is set at 2\\% - 3\\%.  The 2\\% model has a ($V$-weighted) mean age of 11.7 Gyr, with an abundance distribution FWHM scatter of .3 dex.  This is the only model with a mass spectrum (as opposed to the $F_{ini}$ plus constant $dM$ models) that fits all of the observational constraints. Consequently, even though this constant power spectrum model is not physical, the good fit seems to indicate that there will be {\\em some} power spectrum that will satisfy the constraints equally well.  Other variants were tried. We tried superimposing two different schemes atop one another. For instance, suppose half of the galaxies were truly ancient and the other half was allowed to have mergers. These superposition schemes fail dramatically. Unless one is exquisitely artful one ends up with Mg-$\\sigma$ diagrams that are clearly double.  Finally, we tried schemes of variable gas fraction, where the early universe was assumed to be gas rich, but later mergers are mostly stellar. Similarly to the regular 4-parameter models, envelope parameters $n>-0.5$ are ruled out, but uniform-with-time probability gives good matches for $20 < N_{merger} < 40$. In fact, one such model, with F$_{gas}$ varying linearly with time from 0.9 at $z_f$ to 0.1 today matches the Mg-$\\sigma$ data best of all models. This model is also intriguing because it shows strong Mg-$\\sigma$ asymmetry evolution at only modest redshift (see next section).  \\subsection{Predictions for Lookback Studies}  The recent proliferation of 8-12m telescopes will greatly accelerate the rate at which Mg-$\\sigma$ data will become available for distant galaxies. The result of studies to date, that the drift of the Mg zero point is ``consistent with passive evolution,'' \\citep{Zieg97} will eventually be broadened to study the shape of Mg residuals as well. We therefore present some model predictions for high-redshift studies. There is a major caveat to consider here: model galaxies are assumed to be ``star forming'' for 50 Myr after a burst, but ``early type'' at all other times. If merging is a dominant process, then premerger fragments are often likely to be spiral galaxies. If semianalytic models like \\citet{kauf92} are correct, then, over time, some ellipticals can accrete a gas disk and become spirals. So the predictions we show in this section are of questionable validity for merger-dominated models (i.e. all of those that match the $z=0$ data!), but nevertheless provide a useful first cut for the variety of behavior we might expect for Mg-$\\sigma$ relations at significant lookback times.  \\placefigure{fig8}  Figure \\ref{fig8} shows a collection of near-best models as they would appear if observed at large lookback times. The merging trees of the model galaxies are computed to redshift zero, but only the partially complete history of each galaxy is considered for nonzero redshift. So the further back in time we look, the more incomplete the galaxies are. Figure \\ref{fig8} at a glance shows that Mg-$\\sigma$ relations survive intact even under fairly severe merging scenarios back to when the Universe was a quarter of its present age.  The bottom panel shows Mg$_{300}$ as a function of redshift. Within a zeropoint shift, all models show Mg evolution very similar to the ``ancient, passive evolution'' case shown in large diamond symbols at strong Mg strength. It is easy to shift any given model track by adjusting the assumed metallicity scale, so the overall Mg level by itself is not a strong constraint, but the relative time evolution should be fairly well modeled. The most shallow slope, that is, the most constant Mg level with time, is given by the $F_{gas}=0.25$, 40-merger, uniform probability, $F_{ini}=0.5$ model (small open diamonds). This is a merger-heavy model. We think that some researchers will be surprised by the fact that a merger-driven model is redder, relatively speaking, at $z=1.5$, than pure passive evolution scenario because one usually thinks of mergers as causing more star formation activity and hence bluer colors and weaker line strengths. In reality, it can go either way depending on the model. The steepest Mg track is from a very similar model but with the variable $F_{gas}$ option and $F_{ini}=0$ (filled bullets). Neither steepest nor shallowest models differ by more than 0.05 mag in Mg$_{300}$ from the passive evolution case if they are normalized to match at one end or the other of the illustrated redshift range.  The tightness and asymmetry of the Mg-$\\sigma$ relation also do not show large changes with lookback time. Both the ADR and ADev can increase or decrease according to the model, but at a very modest level that will be hard to measure without a very large sample of distant galaxies."
    },
    "0206/quant-ph0206034_arXiv.txt": {
        "abstract": "We discuss an experiment conducted by Nesvizhevsky et al. As it is the first experiment claimed to have observed gravitational quantum states, it is imperative to investigate all alternative explanations of the result. In a student project course in applied quantum mechanics, we consider the possibility of quantummechanical effects arising from the geometry of the experimental setup, due to the \"cavity\" formed. We try to reproduce the experimental result using geometrical arguments only. Due to the influence of several unknown parameters our result is still inconclusive. ",
        "introduction": "A wellknown property of quantum mechanics is the quantisation of the energy levels of a particle trapped in a potential well. For instance, the electromagnetic and the strong nuclear force create different kinds of potential wells and are responsible for many observed phenomena in nature, such as the structure of atoms and nuclei. This suggests that a splitting of the energy levels should also be observed for particles in the Earth's gravitational field. But, since the gravitational field is much weaker, the effect should be subtle and hard to detect.  In a letter to Nature \\cite{Nev1}, Nesvizhevsky et al. claim to have observed such quantum effects of gravity acting on ultracold neutrons (UCNs). They conducted an experiment in which the UCNs were allowed to flow through a cavity with a reflecting surface below and an absorber above. By measuring the number of neutrons exiting the experimental setup, they claim to have observed discrete energy levels. However, in their argument they seemingly disregard the modification due to the absorber, stating that it is ``sufficiently perfect''. Further they argue that the discretisation is related to the sudden increase of neutrons coming through at distinct widths between the reflecting surface and absorber. However, since the UCNs are restricted by both the reflecting surface and the absorber, also the geometrical effects should be considered. The results might even be explained by geometrical arguments only.  The aim with this report is to show that by only using geometrical arguments, the effects observed by the scientist at Grenoble can be explained. ",
        "conclusions": "The main problem in our attempts to explain the results with a geometrical model is that $N_{max}(z)$ is unknown, depending on the spread of energies in the neutron beam. Other difficulties are finding realistic potentials describing the absorber and that, due to a lack of data, our treatment is static. The time independence will exclude important phenomena such as tunnelling effects.  Hence our results are inconclusive. We have not yet determined a potential satisfying our requirements and we await the next report from the experimental group, hoping it will contain the information we need.  We propose the following improvements of the experiment: \\begin{itemize} \\item Rotating the experimental setup by $90^\\circ$ keeping everything else, especially the transverse neutron energies, constant. If the same result occurs it would indicate that the result is due only to the geometry of the experimental setup as no gravitational quantum states can form in this case. \\item Increasing the length of the cavity. If the output would decrease it might confirm our theory of an absorption per unit length. \\item A measurement of where the neutrons strike the detector. The probability distribution should reflect $|\\psi|^2$ (since there is a standing neutron wave, the neutron is not falling in a classical sense). \\end{itemize}  \\begin{appendix}"
    },
    "0206/astro-ph0206285_arXiv.txt": {
        "abstract": "A detailed hydrodynamical model of the gas flow in the triaxial gravitational potential of the bulge of the Andromeda galaxy (M31) has recently been proposed by Berman (2001), and shown to provide excellent agreement with the CO emission line velocities observed along its major axis. In the present paper, we confirm the validity of that model by showing that it can also reproduce the CO velocities observed off the major axis -- a much more robust test. The CO observations, however, tend to span a wider range of velocities than a direct application of the original model of Berman would suggest. This situation can be improved significantly if the molecular disk is made thicker, a requirement already encountered in dynamical simulations of other spiral galaxies, and typically attributed to a broadening of the molecular layer in galactic fountain--like processes. In the central regions of M31, however, it is unclear whether there actually is a thick molecular disk, or whether broadening the molecular layer is merely an artificial theoretical means of accounting for some disk warping. Other effects not included in the model, such as hydraulic jumps, might also contribute to a widening of the velocities. ",
        "introduction": "There is growing evidence that spiral galaxies in which no bar is easily seen on optical or infrared images, nonetheless, have triaxial bulges. Twisting of the inner isophotes and misalignments between the disk and the bulge major axes (both of which have been detected in several spirals with no obvious bar) are clear signposts of triaxiality. Even stronger cases have been made by combining optical or infrared photometric data with spectroscopic observations of the gaseous interstellar component. Indeed, when submitted to a triaxial gravitational potential, the interstellar gas near the centre of a spiral galaxy can be found at higher velocities than expected from circular motion. Such anomalous velocities have been reported in several galaxies where no strong bar is visible, including the two nearest examples: the Milky Way (Rougoor \\& Oort 1959; Dame et al.\\ 2001) and the Andromeda galaxy (Lindblad 1956; Loinard et al.\\ 1999).  The triaxiality of spiral galaxy bulges is particularly important for two reasons. First, it can affect the dynamics of interstellar gas, gathering large amounts of it near galactic centres. This could help to explain the appearance of active galactic nuclei. On long timescales, a build of gas in galactic centres could affect the morphology of spiral galaxies, and imply secular changes in their Hubble types. Second, triaxiality allows more constraints to be derived from spectral observations. In an axisymmetric system, the rotation curve only allows the determination of the radial mass distribution, without providing any three dimensional information. In particular, it is not possible to assess in an axisymmetric system whether the dark matter component is distributed in a flat disk-like structure or in a spherical halo. In a triaxial system, that disk--halo degeneracy can be removed, with obviously important consequences for our understanding of dark matter. Cold Dark Matter (CDM) models, for instance, favor spherical halos.  Since triaxiality primarily affects the dynamics at the centres of spiral galaxies, it should be sought in our neighbours, where the achievable linear resolution is highest. This is particularly important when using radio observations of the gaseous component, which often have limited angular resolutions. Detailed studies of the dynamics of the central regions of the Milky Way are adversely affected by our location inside of the system, and the well-known associated confusion and distance ambiguity problems. However, the existence of triaxial features in the inner Galaxy have been postulated for over forty years. Based on the discovery of the 3 kpc arm, Kerr (1967) proposed the existence of a gaseous bar at the Galaxy centre, and Mulder \\& Liem (1986) modelled the arm as a density wave in a barred potential. By combining photometric and kinematic data, Gerhard \\& Vietri (1986) demonstrated the triaxiality of the Galactic bulge, and COBE images of the Galactic centre have confirmed this triaxiality (Dwek et al.\\ 1995; Binney et al.\\ 1997).  The next nearest spiral (the Andromeda Galaxy -- M31) has no clear visible bar, and is an excellent target for a dynamics study, because its high inclination enables an accurate determination of the gas kinematics. M31 has long been known to have twisted inner isophotes and misaligned bulge and disk major axes (Lindblad 1956). Moreover, the existence of anomalous gas velocities in the inner few kiloparsecs is also now well established both for the atomic (Brinks \\& Burton 1984) and molecular (Loinard et al.\\ 1996, 1999) components. A convincing interpretation in terms of triaxiality was put forward as early as 1956 by Lindblad, and was further developed by Stark (1977) and Stark \\& Binney (1994). However, those early works were hampered by the lack of strong observational constraints.  The distribution and kinematics of the atomic component of the interstellar medium (ISM) has long been known at high resolution and sensitivity across the entire disk of M31 thanks to interferometric observations of the 21--cm line of atomic hydrogen (Brinks \\& Shane 1984, and references therein). However, those observations do not provide clear information on the ISM kinematics in central regions, because the kinematic component associated with the warped outer disk can be seen in projection through the inner disk. The resulting confusion between inner and outer disk seen in the \\HI\\ position-velocity diagrams of the central regions of M31 precludes accurate studies of its dynamics. The molecular component traced by CO emission is a better tracer of the dynamics of the inner regions because no CO emission can be detected in the warped outer disk. However, the inner regions of M31 are also particularly dim in CO (Dame et al.\\ 1993, Loinard et al.\\ 1999). Fragmentary CO observations of the central regions of M31 were used to constrain the model of Stark \\& Binney (1994). Recently, however, more systematic CO observations have been obtained.  Loinard et al.\\ (1995) presented a deep search for CO emission along the inner major axis of M31. This search confirmed the existence of anomalous velocities in the inner few kiloparsecs of M31, and results along the line of nodes of the disk were used by Berman (2001) to constrain an improved dynamical model of the central regions of M31 (see \\S 2 below). However, Berman (2001) pointed out that off-axis CO observations would provide better constraints. Such data are now available -- at least for the southern half of M31 -- thanks to the CO(1-0) survey made at the {\\em Five College Radio Astronomy Observatory} (FCRAO -- Loinard et al.\\ 1996, 1999). The angular resolution of the survey is $\\sim$ 1 arcmin, and it covers the entire Southern part of M31 with a sampling of 50 arcsec. The observational noise level is rather constant across the entire surveyed region with a typical r.m.s. of 45--50 mK per 3.25 km s$^{-1}$ spectral channel. Although the CO(1-0) is found to peak in the broad Population I ring at 10 kpc from the centre, significant emission is still detected in the inner 2--3 kpc (10--15 arcmin), where most constraints can be obtained about the triaxiality of the bulge. The integrated intensity images or position--velocity diagrams shown in Loinard et al.\\ (1999), as well as the position--velocity diagrams that will be shown in this article, have been obtained after `unsharp masking' was applied to the data. In this scheme, a data cube smoothed both in position and velocity (but not regridded) is first constructed from the original data. The average noise level $\\sigma$ in this smoothed cube is computed, and all the pixels in the non--smoothed data cube where no significant ($>$ 3$\\sigma$) emission is found in the smoothed version are blanked. This avoids that unnecessary noise is added when summations over several pixels are performed, and in effect reduces the noise in integrated maps by a factor of a few. A slight drawback of this method is that the noise becomes dependent on the number of pixels summed (or averaged), and, therefore, on the position in the maps.  Anomalous velocities are clearly seen in this data set in the inner regions of M31 (see fig.\\ 12 in Loinard et al.\\ 1999). While the cut along the major axis essentially shows the same anomalous velocities as reported by Loinard et al.\\ (1995), cuts parallel to the major axis provide new constraints for the models, as called for by Berman (2001). ",
        "conclusions": "In this article, we show that hydrodynamical triaxial models of the bulge of the Andromeda galaxy (M31) previously presented by Berman (2001), and initially tested against CO emission observations along the apparent major axis only, can also account for the CO emission velocities observed off the apparent major axis -- a much more robust validity test. The triaxial model implies that the bulge of M31 is a fast rotator and hence the dark matter contained in its central regions must be minimal and coincident with the stellar disk.  To account for the whole velocity range covered by the observations, a finite, fairly large thickness (half width at half maximum intensity of 200--500 pc) must be given to the molecular disk. It is quite unclear, however, whether the inner molecular disk of M31 is truly that thick or whether adding it to the model is merely an effective theoretical trick able to account for a warping of a thiner disk. Other out of plane, or peculiar velocities (such as those associated with hydraulic jumps or turbulence) could also contribute to the wide velocity range covered by the CO observations."
    },
    "0206/astro-ph0206370_arXiv.txt": {
        "abstract": "We present the optical, near-infrared, submillimeter, and radio follow-up catalog of the X-ray selected sources from an $\\approx 1$~Ms {\\it Chandra} observation of the Hubble Deep Field North region. We have $B$, $V$, $R$, $I$, and $z'$ magnitudes for the 370 X-ray point sources, $HK^\\prime$ magnitudes for 276, and spectroscopic redshifts for 182. We present high-quality spectra for 175 of these. The redshift distribution shows indications of structures at $z=0.843$ and $z=1.0175$ (also detected in optical surveys) which could account for a part of the field-to-field variation seen in the X-ray number counts; however, these structures do not dominate the number of X-ray sources in the sample and hence should not strongly affect the redshift distribution. All of the X-ray sources with $z>1.6$ are either broad-line AGN or have narrow Ly$\\alpha$ and/or CIII]~1909~\\AA\\ emission; none of the known $z>1.6$ absorption-line galaxies in the field are detected individually in X-rays. We estimate photometric redshifts for the sources with $B-I>1.5$ (bluer than this it is hard to distinguish between low redshift irregulars and luminous high-redshift AGN) and find agreement (most are within 25\\%) with the available spectroscopic redshifts. The majority of the galaxies in both the $2-8$~keV (hard) and $0.5-2$~keV (soft) samples have absolute magnitudes comparable to or more luminous than $M_I^\\ast=-22$. The flux contributions separated into unit bins of redshift show that the $z<1$ spectroscopically identified sources already contribute about one-third of the total flux in both the hard and soft bands. Thus, major accretion onto supermassive black holes has occurred since the Universe was half its present age. We find from ratios of the X-ray counts that the X-ray spectra are well-described by absorption of an intrinsic $\\Gamma=1.8$ power-law, with $N_H$ values ranging from about $10^{21}$~cm$^{-2}$ to $5\\times 10^{23}$~cm$^{-2}$. We find very little evolution in the maximum rest-frame opacity-corrected and $K$-corrected $2-8$~keV X-ray luminosities with decreasing redshift until $z\\lesssim 0.5$, where the volume becomes too small to probe effectively very high luminosity sources. We estimate that the {\\it Chandra} sources that produce 87\\% of the {\\it HEAO-A} X-ray background (XRB) at 3~keV produce 57\\% at 20~keV, provided that at high energies the spectral shape of the sources continues to be well-described by a $\\Gamma=1.8$ power-law. However, when the {\\it Chandra} contributions are renormalized to the {\\it BeppoSAX} XRB at 3~keV, the shape matches fairly well the observed XRB at both energies.  Thus, whether a substantial population of as-yet undetected Compton-thick sources, or some change in the spectral shape of the current sources from the simple power-law dependence, is required to completely resolve the XRB above $10$~keV depends critically on how the currently discrepant XRB measurements in the $1-10$~keV energy range tie together with the higher energy XRB. ",
        "introduction": "\\label{secintro}  X-ray surveys most directly trace accretion onto supermassive black holes and hence provide our best window on black hole evolution. Some primary observational goals of X-ray studies include the measurement of supermassive black hole properties (such as number density and accretion rates) versus redshift and the determination of the nature of the host galaxies. Since much of the accretion power in the Universe may be absorbed by substantial neutral hydrogen column densities (e.g., \\markcite{fabian99}Fabian \\& Iwasawa 1999), the challenge is to construct a {\\it complete census} of supermassive black holes to the earliest epoch, including sources which are heavily obscured from soft X-ray energies to the near-infrared (NIR). The census is well underway for unobscured sources (i.e., quasars with broad emission lines and big blue bumps), but the obscured population is not well-characterized, except for the relatively rare radio galaxies.  A fundamental goal of the {\\it Chandra X-ray Observatory} was to resolve the hard ($2-8$~keV) X-ray background (XRB) into discrete sources. At these energies the photons can penetrate all but the highest column densities ($>10^{24}$~cm$^{-2}$) of gas and dust, so many obscured active galactic nuclei (AGN) can now be detected. The two recent $\\approx 1$~Ms exposures of the {\\it Chandra} Deep Field North (CDF-N; \\markcite{brandt01}Brandt et al.\\ 2001c, hereafter B01) and the {\\it Chandra} Deep Field South (CDF-S; \\markcite{giacconi02}Giacconi et al.\\ 2002) have resolved $>80-90$\\% of the $2-8$~keV XRB into discrete sources (\\markcite{campana01}Campana et al.\\ 2001; \\markcite{cowie02}Cowie et al.\\ 2002; \\markcite{rosati02}Rosati et al.\\ 2002) with the main uncertainty being the normalization of the XRB itself. The CDF-N exposure has recently been extended to a second megasecond.  Although with the {\\it Chandra} and {\\it XMM-Newton Observatories} there has been rapid improvement in our ability to detect hard X-ray sources at the faintest fluxes, our efforts to understand in detail the nature and evolution of the sources creating the hard XRB are still in the early stages. Optical and NIR follow-up studies of $100-300$~ks {\\it Chandra} (\\markcite{mushotzky00}Mushotzky et al.\\ 2000; \\markcite{horn01}Hornschemeier et al.\\ 2001, hereafter H01; \\markcite{barger01a}Barger et al.\\ 2001a, c; \\markcite{tozzi02}Tozzi et al.\\ 2002; \\markcite{stern02b}Stern et al.\\ 2002b) and {\\it XMM-Newton} (\\markcite{hasinger01}Hasinger et al.\\ 2001; \\markcite{baldi02}Baldi et al.\\ 2002; \\markcite{fiore01}Fiore et al.\\ 2001) imaging surveys have found that almost half of the hard X-ray light arises in optically bright ($I<23.5$) galaxies in the $z<1.5$ redshift range. Almost all of these sources can be spectroscopically identified, and most are bulge-dominated galaxies with near-$L^\\ast$ luminosities. The other half of the hard X-ray light arises in a mixture of $z>1.5$ AGN and optically faint galaxies ($I>23.5$) that cannot be spectroscopically identified.  Contrary to the situation for the faint {\\it ROSAT} soft X-ray sources (\\markcite{hasinger98}Hasinger et al.\\ 1998; \\markcite{schmidt98}Schmidt et al.\\ 1998), the vast majority ($>80$\\%) of the spectroscopically identified hard X-ray sources do not have broad optical or ultraviolet lines, and almost half show no obvious high-ionization signatures of AGN activity in their spectra. The hard to soft X-ray flux ratios of the latter sources suggest that the sources are highly absorbed systems whose high column densities could effectively extinguish the optical, ultraviolet, and NIR continua from the AGN and render traditional identification techniques ineffective.  Based on the colors, many of the spectroscopically unidentified $I>23.5$ galaxies may lie at redshifts in the range $z=1.5$ to 3 (\\markcite{crawford01}Crawford et al.\\ 2001; \\markcite{cowie01}Cowie et al.\\ 2001; \\markcite{barger01a}Barger et al.\\ 2001a, b; \\markcite{alexander01}Alexander et al.\\ 2001). However, a very intriguing possibility suggested by \\markcite{haiman99}Haiman \\& Loeb (1999) is that some of the optically faint sources may instead be low luminosity quasars at very high redshifts ($z>5$), since cold dark matter models can allow a large number of such high-redshift AGN.  The $\\approx 1$~Ms {\\it Chandra} exposures have also resolved $>90$\\% of the soft ($0.5-2$) XRB into discrete sources (B01; \\markcite{rosati02}Rosati et al.\\ 2002). At the faintest X-ray fluxes probed by these exposures, galaxies whose X-ray light is dominated by processes related to star formation rather than AGN activity begin to be significant contributors to the number counts, even though they contribute only $2-3$\\% to the X-ray light (e.g., \\markcite{brandt01a}Brandt et al.\\ 2001a; \\markcite{horn02}Hornschemeier et al.\\ 2002). The large numbers of detections are understandable since we now have the sensitivity to detect the most luminous ($>2\\times 10^{39}$~erg~s$^{-1}$) X-ray binaries and supernova remnants at $z<0.15$. AGN may cease to dominate the X-ray number counts at extremely faint fluxes (e.g., \\markcite{moran99}Moran, Lehnert, \\& Helfand 1999) in a manner analogous to the dominance of star forming galaxies rather than AGN at very faint radio fluxes (e.g., \\markcite{richards00}Richards\\ 2000).  In this paper we present follow-up optical, NIR, and submillimeter imaging and optical spectroscopy of the X-ray sources detected in the $\\approx 1$~Ms CDF-N exposure to characterize the properties of the faint X-ray source populations. We then use our data to describe the evolution of the X-ray sources and their host galaxies. ",
        "conclusions": "\\label{secsummary}  We presented a catalog of the optical, NIR, submillimeter, and radio properties of the X-ray sources identified in the $\\approx 1$~Ms {\\it Chandra} exposure of the CDF-N. We now have redshifts for 182 of the 370 X-ray sources, and we presented the spectra for 175. The redshift identifications are very complete (78\\%) for the $R\\le 24$ galaxy sources in a $10'$ radius around the approximate X-ray image center. All of the X-ray sources with $z>1.6$ are either broad-line AGN or have narrow Ly$\\alpha$ and/or CIII]~1909~\\AA\\ emission; none of the known $z>1.6$ absorption line systems in the field are detected in X-rays.  We found spectroscopic evidence for large scale structure in the field which could account for a part of the field-to-field variation seen in the X-ray number counts. However, since the observed structures do not dominate the number of X-ray sources in the sample, the redshift distribution should not be too strongly affected by clustering.  The broad-line sources are all extremely blue, making it hard to distinguish between low redshift irregulars and luminous high-redshift AGN purely on the basis of color. Very few of the redder X-ray sources, whether spectroscopically identified or not, lie outside the regions populated by unevolving spiral galaxy tracks.  We estimated photometric redshifts for the sources using our five broad-band colors (six when $HK'$ was available). Above $B-I=1.5$ the photometric redshift estimates are quite robust; for most sources with available spectroscopic redshifts the photometric redshifts are within 25\\%. The majority of galaxies have absolute magnitudes comparable to or more luminous than $M_I^\\ast=-22$.  We found that the spectral energy distributions of the X-ray sources are well-described by a $\\Gamma=1.8$ intrinsic spectrum corrected for absorption. We determined the absorbing column densities for our sources with redshifts and found that they range from about $10^{21}$~cm$^{-2}$ to $5\\times 10^{23}$~cm$^{-2}$.  We calculated intrinsic rest-frame hard and observed-frame soft X-ray luminosities. We found very little evolution in the maximum hard X-ray luminosities until $z\\lesssim 0.5$, at which point the volume becomes too small to probe the very high luminosity sources. At soft X-ray luminosities of $L_{0.5-2~keV}<10^{42}$~erg~s$^{-1}$ the `typical' galaxy, whose X-ray emission may be dominated by a population of X-ray binaries, hot interstellar gas, or a low-luminosity AGN, becomes important.  We calculated $B$-band luminosities, $L_B$, for the X-ray sources. Sources with $z>1.2$ approximately follow the line $\\log(L_B/L_{2-8~keV})=1$, which suggests that the AGN dominate the optical luminosities. At $L_{2-8~keV}\\sim 10^{43}$~erg~s$^{-1}$ the locus flattens above the line, which suggests this is where the host galaxy light starts to dominate the optical light.  Redshift slices of the sources versus X-ray flux show that substantial contributions to the total X-ray flux in both bands (of order one-third using only the present spectroscopic sample) come from $z<1$. Thus, major accretion onto supermassive black holes has occurred since the Universe was half its present age.  We estimated that the {\\it Chandra} sources that produce 87\\% of the {\\it HEAO-A} XRB at 3~keV produce 57\\% at 20~keV, provided that at high energies the spectral shape of the sources continues to be well-described by a $\\Gamma=1.8$ power-law. However, when the {\\it Chandra} contributions are renormalized to the {\\it BeppoSAX} XRB at 3~keV, the shape matches fairly well the observed XRB at both energies. Thus, whether a substantial population of as-yet undetected Compton-thick sources or some change in the spectral shape of the current sources from the simple power-law dependence is required to resolve the high energy XRB depends critically on how the low energy and high energy XRB measurements tie together."
    },
    "0206/astro-ph0206056_arXiv.txt": {
        "abstract": "We test the effectiveness of photometric redshifts based upon galaxy spectral template fitting for X-ray luminous objects, using a sample of 65 sources detected by \\chandra in the field of the Caltech Faint Galaxy Redshift Survey (CFGRS).  We find that sources with quasar-dominated spectra (for which galaxy spectral templates are not appropriate) are easily identified, and that photometric redshifts are robust for the rest of the sources in our sample. Specifically, for the 59 sources that are not quasar-dominated at optical wavelengths, we find that the photometric redshift estimates have scatter comparable to the field galaxy population in this region.  There is no evidence for a trend of increasing dispersion with X-ray luminosity over the range $L_X=10^{39} - 5\\times10^{43}$ erg s$^{-1}$, nor is there a trend with the ratio of X-ray to optical flux, $f_X/f_R$. The practical implication of this work is that photometric redshifts should be robust for the majority ($\\sim$90\\%) of the X-ray sources down to $f_X\\approx10^{-16}$ erg s$^{-1}$ cm$^{-2}$ that have optical counterparts brighter than $R\\approx24$. Furthermore, the same photometry can be easily used to identify the sources for which the photometric redshifts are likely to fail.  Photometric redshift estimation can thus be utilized as an efficient tool in analyzing the statistical properties of upcoming large \\chandra and \\xmm data sets and identifying interesting subsamples for further study. ",
        "introduction": "\\label{sec:introduction} Photometric redshift estimation is a powerful tool in extragalactic astronomy, and significant effort has been expended in recent years developing robust redshift estimation techniques. The most widely employed approach is to use spectral templates (either empirical or from stellar synthesis models) to derive optimal fits to the observed galaxy colors \\citep[e.g.][]{lanzetta1996, fernandezsoto1999, benitez2000,furusawa2000}.  This approach has been highly successful for normal galaxies, achieving results as good as $\\sigma_z=0.06(1+z)$ for the {\\it Hubble} Deep Field \\citep[HDF;][]{fernandezsoto1999, benitez2000, furusawa2000}.  Extending upon this work, \\citet{budavari2001} and \\citet{richards2001} have recently found that photometric redshifts accurate to within $\\Delta z=0.2$ can be obtained for 70\\% of quasars in the Sloan Digital Sky Survey -- even at $z<2.2$, where quasar spectra lack a strong continuum break in the Sloan $ugriz$ filters.  In this work we focus on X-ray luminous objects, asking whether a self-consistent technique can be devised to obtain photometric redshifts for X-ray selected samples.  Recent work on deep \\chandra fields indicates that the resolved X-ray background is comprised of a variety of objects, including early-type ellipticals, starburst galaxies, obscured active galactic nuclei (AGN), narrow- and broad-line AGN, and quasars \\citep[see e.g.][]{mush2000,barger2001,horn2001,tozzi2001,stern2002,barger2002}. Given the disparate nature of these sources, the effectiveness of photometric redshifts is {\\it a priori} unclear. Recent spectroscopic observations in the \\chandra Deep Field-North (CDF-N) highlight this concern.  \\citet{barger2002} find that, out of a sample of 182 hard sources (2-8 keV detections) with spectroscopic follow-up, approximately half show signatures of AGN activity in their spectra.   As argued by \\citet{barger2002}, one can reasonably expect that traditional galaxy spectral template fitting should be reliable for the 50\\% of sources that lack an AGN spectral signature.  Similarly, it is likely that quasar spectral templates can be utilized for AGN-dominated sources (although this expectation has not yet been verified). It is unclear though whether either approach is robust for the intermediate sources in which both the AGN and the galaxy stellar population contribute significantly to the optical spectrum. Do photometric redshifts based upon galaxy spectral template fitting gradually degrade with increasing fractional luminosity contribution from the AGN, or do they remain robust until the AGN contribution reaches some critical level? The \\hubble Deep Field, which is the canonical field for testing photometric redshifts, offers little insight. There are only six X-ray luminous sources with spectroscopic redshifts in the HDF -- the majority of which are low-luminosity. Photometric redshift comparisons in the HDF, such as the blind check of \\citet{cohen2000}, thus cannot test the reliability of photometric redshift estimators in the interesting regime of properties.\\footnote{ One of these sources is an outlier in the blind photometric redshift tests of \\citet{cohen2000}; however, it has also been argued by \\citet{fernandezsoto2001} that the spectroscopic redshift for this source is incorrect. Two of the other objects are listed as faint X-ray sources in \\citet{horn2001}, but do not appear in the $\\approx$1 Ms \\chandra catalogs of \\citet{brandt2001}.}  There are strong practical motivations for using an expanded sample of X-ray selected sources to assess the reliability of photometric redshifts for these objects.  Foremost, an efficient means of redshift estimation is required to maximize the return from large area \\chandra and \\xmm surveys, such as the Lockman Hole survey ({\\it Chandra}, PI: Barger), the \\chandra Multiwavelength Project \\citep[ChaMP][]{wilkes2000}, the {\\it XMM-}LSS survey \\citep{pierre2001}, and the upcoming \\chandra survey in the NOAO Deep-Wide Field (9 sq. degrees, PIs: Jones \\& Murray). A number of the issues that these surveys aim to address, such as evolution in the AGN X-ray luminosity  \\begin{inlinefigure} \\plotone{f1.eps} \\figcaption{ Optical $R$-band magnitude from \\citet{barger2002} plotted against the full band 0.5-8 keV X-ray flux from \\citet{brandt2001}.  The crosses correspond to the seven sources in our sample with colors that are best fit by a quasar spectral template.  The filled circles correspond to the other sources in our sample, while the open circles denote the rest of the extragalactic sources in the \\citeauthor{barger2002} catalog that lie within the CFGRS region. Arrows denote sources that only have lower (upper) limits on their $R$-band magnitudes (X-ray fluxes). The dashed lines denote lines of constant $f_X/f_R$, as given by Equation 3 in \\citet{horn2001}. \\label{fig:scatter}} \\end{inlinefigure}   \\noindent function and evolution in the correlation function of faint X-ray sources, are statistical in nature and do not require the accuracy of spectroscopic redshifts.  For example, the survey in the NOAO Deep Wide Field aims to study the X-ray spatial correlation function using an expected $\\sim$2200 sources (C. Jones 2002, private communication). This project has no associated large spectroscopic program and will depend upon the use of photometric redshifts to achieve its key scientific goal.  If one can attain a level of accuracy for these photometric redshifts comparable to what is achieved for ``normal'' galaxies, photometric redshifts will be a valuable tool. If not, then the limitations of photometric redshift techniques must be established.  In this paper we assess the validity of photometric redshifts in the Caltech Faint Galaxy Redshift Survey field \\citep[CFGRS,][]{cohen2000,hogg2000} using the \\citet{brandt2001} catalog of X-ray sources in the CDF-N. We first compare the scatter in photometric redshift estimates for X-ray selected sources with a control sample of quiescent galaxies, and quantify the fraction of X-ray sources with bright optical counterparts for which traditional galaxy spectral template fitting is valid. Next, we search for trends between the redshift residuals and other physical quantities, testing whether there is a range in X-ray luminosity ($L_X$) or flux ratio ($f_X/f_R$) over which photometric redshifts degrade. Finally, we ask whether the objects for which photometric redshifts are likely to fail (e.g. quasars) can be easily identified using the optical photometry. The data are described in \\S \\ref{sec:data} and the results are presented in \\S \\ref{sec:photoz}.  In \\S \\ref{sec:summary} we summarize our work and discuss prospects for future studies with larger and fainter samples. We assume $\\Omega_0=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=100 h$ km s$^{-1}$ Mpc$^{-1}$.  \\begin{inlinefigure} \\plotone{f2.eps} \\figcaption{ Template quasar spectra used in the photometric redshift code. The dashed line separates the SDSS composite spectra (short wavelength), which is corrected for stellar contamination, from the PDS 456 spectra.  \\citet{simpson1999} obtained $J-$ and $K-$band spectra for PDS 456; between these two bands we interpolate. \\label{fig:qsotemplate}} \\end{inlinefigure} ",
        "conclusions": "\\label{sec:summary}  We use a sample of X-ray sources detected by \\chandra in the Caltech Faint Galaxy Redshift Survey region to test the robustness of photometric redshifts for these objects. For the 59 out of 65 sources with colors that are not best fit by a quasar spectral template, we find no degradation in the accuracy of photometric redshifts as compared to the field galaxy population. We also find that the redshift residuals exhibit no trend as a function of X-ray luminosity or $f_X/f_R$ for these sources.  We demonstrate that it is feasible to quickly identify objects whose spectra are quasar-dominated at optical wavelengths and derive robust photometric redshifts for the other $\\sim90$\\% of sources that have optical counterparts brighter that $R\\approx24.5$, which includes roughly two thirds of all sources with $f_X>10^{-16}$ erg s$^{-1}$ cm$^{-2}$ \\citep{barger2002}. Consequently, photometric redshift estimation should be a valuable tool for upcoming large samples of optically bright, X-ray selected objects.  This paper is a first step towards quantifying the robustness of photometric redshifts for systems in the transition regime between quiescent galaxies and quasars. The next required step is extension to fainter optical magnitudes.  Given that the majority of fainter sources have high $f_X/f_R$, it is unclear whether photometric redshifts will be as effective for this population, but the motivation for testing their robustness is strong. There is speculation that these optically faint, high $f_X/f_R$ sources possibly include obscured AGN at $z\\ga5$ \\citep[e.g.][]{stern2002}, and photometric redshifts may be the only viable approach for studying this population."
    },
    "0206/astro-ph0206260_arXiv.txt": {
        "abstract": "We numerically integrated the orbits of 1458 particles in the region of the classical Kuiper Belt (41 AU $\\leq a \\leq$ 47 AU) to explore the role of dynamical instabilities in sculpting the inclination distribution of the classical Kuiper Belt Objects (KBOs).  We find that the selective removal of low-inclination objects by overlapping secular resonances ($\\nu_{17}$ and $\\nu_{18}$) acts to raise the mean inclination of the surviving population of particles over 4 billion years of interactions with Jupiter, Saturn, Uranus and Neptune, though these long-term dynamical effects do not themselves appear to explain the discovery of KBOs with inclinations near $30^{\\circ}$.  Our integrations also imply that after 3 billion years of interaction with the massive planets, high inclination KBOs more efficiently supply Neptune-encountering objects, the likely progenitors of short-period comets, Centaurs, and scattered KBOs.  The secular resonances at low inclinations may indirectly cause this effect by weeding out objects unprotected by mean motion resonances during the first 3 billion years. ",
        "introduction": "Most of the mass of the Kuiper Belt between 30 and 50 AU appears to reside in a region outside semimajor axis $a = 40$ AU called the classical Kuiper Belt \\citep{jewi98}.  Early Kuiper Belt surveys found that Kuiper Belt Objects (KBOs) can have surprisingly high orbital inclinations \\citep{jewi95, jewi96}. Recent large-scale surveys \\citep{jewi98, tjl01} have established that the classical KBOs---not just the relatively nearby plutinos---frequently have high orbital inclinations ($i \\gtrsim 15^{\\circ}$).  The discovery of these high-inclination objects seems to point to unknown processes in the primordial solar system.   Several mechanisms for ``pumping-up'' the ancient KBO inclination distribution have been investigated, such as resonant encounters \\citep{malh95,naga00} and perturbations from passing stars \\citep{ida00}, scattered planets \\citep{thom02} and planetesimals \\citep{peti99}.  Large KBOs that may have existed long ago during the epoch of planet formation could also have stirred the KBO orbits \\citep{keny02}.  However, theories of the ancient solar system alone can not explain the current KBO orbital distribution. Orbit integrations in the Kuiper Belt \\citep{torb89, torb90, glad90, levi93, holm93} have shown that the dynamical effects of the massive planets on the Kuiper Belt have probably removed most of the original KBOs over the lifetime of the solar system.  Objects removed recently by these processes probably supply today's population of Centaurs, short-period comets, and scattered KBOs \\citep{fern80, dunc88}.  One might ask whether selective removal of objects by interactions with the planets has altered the inclination distribution of the surviving population.  Furthermore, debiased estimates of the KBO inclination distribution \\citep{brow01} suggest that the classical KBOs divide into two populations: one dynamically warm, with a typical inclination of $i \\approx 17^{\\circ}$, and one dynamically cold, with a typical inclination $i \\approx 2^{\\circ}$.  We wonder whether long-term dynamical interactions helped create the apparent two-component distribution.  To understand the role of dynamical stability in shaping the classical KBO inclination distribution and creating Centaurs, comets, and scattered KBOs, we performed a new large-scale simulation of the dynamics of KBOs, integrating the orbits of 1458 test particles in the region of the classical Kuiper Belt under the influence of gravity from Jupiter, Saturn, Uranus, and Neptune.  \\citet{dunc95} performed the most recent previous large-scale integration of orbits in the Kuiper Belt.  They used an initial grid with fine resolution in semimajor axis, but only sampled a few initial inclinations, and did not integrate the orbits of high-inclination objects $i > 1^{\\circ}$ for longer than one billion years.  Our work complements theirs; at the cost of coverage in semimajor axis, we focus on exploring the consequences of high inclination and eccentricity over a four-billion-year period. ",
        "conclusions": "Dynamical erosion is not the only long-term effect in the classical Kuiper Belt; collisions may shape the distributions of KBO orbital elements \\citep{ster97, keny98, keny99, durd00}.  But as the collision rate and the mass of the Kuiper Belt decreased with time, the dynamical effects shown by our integrations must have begun to dominate the shaping of the classical Kuiper Belt. We found that interactions with the massive planets preferentially deplete low-inclination objects in the inner half of the classical Kuiper Belt, raising the mean inclination of the population of surviving objects.  However, the outer half of the classical Kuiper Belt ($a > 44$ AU, $q < 39 AU$) could have retained its ancient inclination distribution.   We also found that objects ejected from the classical Kuiper Belt during the last 1 billion years of our integration primarily come from initial orbits at high inclinations and eccentricities. This effect may be caused indirectly by $\\nu_{17}$ and $\\nu_{18}$, which remove many low-inclination particles at early times and weed out those unprotected by MMRs.  This finding implies that Centaurs, short-period comets and scattered KBOs should have an initial color distribution like that of the high-inclination KBOs.  Measuring of the colors of more Centaurs could test this hypothesis.  Perhaps the high-inclination KBOs separated from the primordial KBOs objects during a special event early in the lifetime of the solar system.  If today's active comets arise mainly from high-inclination KBOs, we have this event to thank for their inner solar system company."
    },
    "0206/astro-ph0206110_arXiv.txt": {
        "abstract": "{ In November 2001 we undertook a coordinated observing campaign to study the connection between X-ray and optical variability in the weak-line T Tauri star V410\\,Tau. The observing plan included three $\\sim 15$\\,ksec observations with {\\em Chandra} using the Advanced CCD Imaging Spectrometer for Spectroscopy (ACIS-S) scheduled for different phases of the known 1.87\\,d starspot cycle. Photometric and spectroscopic monitoring of V410\\,Tau involving telescopes on three different continents was scheduled simultaneously with the {\\em Chandra} exposures. ",
        "introduction": "\\label{sect:intro}  V410\\,Tau is a well-known T Tauri star (TTS) of the weak-line class. % Due to its optical brightness ($V \\sim 11$\\,mag) V410\\,Tau has frequently been included in photometric monitoring programs throughout the last three decades. These observations have revealed periodic variability on a $1.87$\\,d cycle, which has remained stable over many years (e.g. Vrba et al. 1988, Herbst 1989, Petrov et al. 1994 hereafter P94). The periodic behavior of V410\\,Tau was attributed to surface spots similar to those observed on BY\\,Dra and RS\\,CVn systems. The amplitude of the asymmetric lightcurve of V410\\,Tau is larger than that of any other weak-line TTS and varies with a time-scale of years between $0.2-0.6$\\,mag in the $V$ band.  Optical photometric and spectroscopic monitoring has occasionally resulted in the detection of flares (Vrba et al. 1988, Guenther \\& Ball 1999). However, in view of the large existing data base for V410\\,Tau, flares are surprisingly rare on this star. This may indicate that the magnetic field of V410\\,Tau maintains a comparatively stable configuration, an idea supported also by the long-term stability of the photometric spot cycle. On the other hand V410\\,Tau is a strong and variable X-ray source (Costa et al. 2000) suggesting that magnetic fields are dissipated in the corona. Despite its obvious variability in the X-ray regime, no direct signs of X-ray flares have been observed on V410\\,Tau so far.  In order to study the connection between the optical and X-ray variability of V410\\,Tau we have undertaken a coordinated observing campaign involving photometric monitoring, optical low- and high-resolution spectroscopy, and {\\em Chandra} observations at different phases throughout the (photometric) spot cycle. Here we report on the first results from the photometric and X-ray monitoring. ",
        "conclusions": "\\label{sect:summary}  The simultaneous X-ray and optical data of this campaign does not support a direct connection between X-ray emitting regions and starspots on V410\\,Tau. In contrast to most previous observations of this star, our continuous monitoring has resulted in the detection of at least one large flare. We tentatively re-determined the rotation period based on the migration of the minimum in the $V$ band lightcurve. Future analysis will show whether this migration could instead be due to a shift of the spot pattern on the surface of V410\\,Tau."
    },
    "0206/astro-ph0206326_arXiv.txt": {
        "abstract": "{ We present the first VLBI images of SiO masers in the circumstellar envelope of an S-type star obtained with the newly developed capabilities of the Very Long Baseline Array (VLBA) at 86 GHz. We combine these data with those obtained quasy simultaneously at 43 GHz.  These observations provide information on the structure and dynamics of the innermost circumstellar shells, where the return of large quantities of stellar material to the interstellar medium starts. Despite of the fact that these are prelimininary results, we report that the maser emission of the $v$=1 J=1--0 and J=2--1 present a very different emission distribution, being not coincident in most cases.  } \\authorrunning{J.-F Desmurs et al.} ",
        "introduction": "SiO masers emission at 7 mm wavelength ($v$=1 and $v$=2, J=1--0 transition near 43 GHz) has been observed in AGB stars with very hight resolution by means of VLBI techniques, yielding important results in relation with their not yet well understood pumping mechanism. The 7 mm maser emission regions are found to be distributed in a number of spots forming a ring-like structure at about 2-3 stellar radii \\citep{col92,dia94,gre95,des00} which are assumed to be centered on the stellar position.  This ring-like flux distribution arise naturally in the framework of the radiative pumping mechanism of SiO masers (see Bujarrabal et al. 1994), but may also be explained by a collisional modelisation.  Recent simultaneous observations of the $v$=1 and $v$=2, J=1--0 (see Desmurs et al. 2000) show that the $v$=1 and $v$=2 maser spots are often close, but appear systematically shifted by a few mas and are only rarely coincident; a result that would argue in favor of radiative pumping models.  In an attempt to pursue the comparison between the predictions of the different models and observational data, we have measured the relative spatial distribution of SiO maser emission between the J=1--0 and J=2--1 transitions. In that case, models using either radiative or collisional pumping predict ``maser chains'' across the excitated vibrational states, in such a way that the inversion of the different-J transition in the same $v$ state are mutually reinforced.  Models predict that the rotational masers in the same vibrational state should appear under the same physical conditions and, therefore, that the 43\\,GHz and 86\\,GHz masers emission originate from the same condensations in the circumstellar envelope (CSE) \\citep[see][]{buj94,hum02}. ",
        "conclusions": "We produced the first VLBA maps with milliarcsecond resolution of the SiO maser emission from an S-type star, $\\chi$~Cyg.  We obtained maps of the transitions $v$=1, J=1--0 at 7\\,mm (see Fig \\ref{chi_cyg_7mm}) and of the $v$=1, J=2--1 at 3\\,mm (see Fig \\ref{chi_cyg_3mm}).  The map size is approximately $\\sim$80 by 80 mas. The gaussian restoring beam has a FWHM of 0.8x0.2\\, mas (with a position angle PA=-12.4$^{\\circ}$) and 0.75x0.05\\, mas (with a position angle PA=-15.8$^{\\circ}$), respectively, for the data at 7 and 3\\,mm. Both figures are using the same scale to ease a direct comparison, but do not share spatial origin as explained before.  Our preliminary results shows that the SiO maser emission distribution in both transition occur more or less at a similar radii from the center.  In particular, we observed that masers for $v$=1, J=2--1 (3\\,mm) arise from a ring-like structure, as it has been reported in several other AGB stars at 43 GHz \\citep{dia94, gre95, des00, phi01}. The radii of this ring structure is of the order of $\\sim$ 28\\,mas which, adopting the Hipparcos distance (ESA 1997) of 106$\\pm$15 pc, is equivalent to 4.5\\,10$^{13}$\\,cm.  One of the most surprising result of these observations is that the emission distribution between the two lines is completly different. Whatever the choosed alignment, it is impossible to make coincident more than one maser spot at the same time. Even the regions emitting in one transition or the other are very different at large scale (see Figs~\\ref{chi_cyg_7mm} and~\\ref{chi_cyg_3mm}). This is in complete contradiction with all theoretical predictions.  \\begin{figure}[t] \\centering \\includegraphics[scale=0.45]{desmurs_fig2.ps} \\vspace{-1cm} \\caption{Integrated intensity maps of the $v$=1, J=2--1 line (rest frequency 86243.442 MHz) of SiO masers towards $\\chi$ Cyg. Contours are multiples by 10\\% of the peak flux ($\\sim$36 Jy). The resolution beam is 0.75x0.052 mas (PA=-15.8$^{\\circ}$). } \\label{chi_cyg_3mm} \\end{figure}  Our conclusions are only preliminary. But in the case that they would be confirmed by future observations and in other sources, this result would cast serious doubts on the present models of the SiO maser excitation."
    },
    "0206/astro-ph0206489_arXiv.txt": {
        "abstract": "We present the first far-ultraviolet (FUV) observations of the magnetic cataclysmic variable VV Puppis, obtained with the {\\em Far Ultraviolet Spectroscopic Explorer} satellite. In addition, we have obtained simultaneous ground-based optical photometric observations of VV Pup during part of the FUV observation. The shapes of the FUV and optical light curves are consistent with each other and with those of past observations at optical, extreme-ultraviolet, and X-ray wavelengths.  Time-resolved FUV spectra during the portion of VV Pup's orbit when the accreting magnetic pole of the white dwarf can be seen show an increasing continuum level as the accretion spot becomes more directly visible. The most prominent features in the spectrum are the \\ion{O}{6} $\\lambda\\lambda$1031.9, 1037.6 emission lines. We interpret the shape and velocity shift of these lines in the context of an origin in the accretion funnel near the white dwarf surface.  A blackbody function with $T_{\\rm bb}\\gtrsim90,000$ K provides an adequate fit to the FUV spectral energy distribution of VV Pup. ",
        "introduction": "VV Puppis was noticed early in the 20th century \\citep{vangent31} as a faint \\cite[$V=14.5$--$18$;][]{downes01}, rapidly periodic ($P\\approx100$ min) variable.  Based on photometric and spectroscopic observations, \\citet{herbig60} suggested it was a binary whose emission lines originate on the brighter component. Later, VV Pup was identified as the third member of the AM Herculis class of magnetic cataclysmic variable \\cite[CV;][]{tapia77}. The white dwarf (WD) primary star in an AM Her system has a magnetic field strength of $B \\gtrsim 10$ MG (currently known to reach up to several hundred MG; e.g., $B = 230$ MG in AR Ursae Majoris, \\citealt{schmidt99}).  The magnetic field prevents the formation of the accretion disk that dominates the luminosity of non-magnetic CVs. Instead, the accretion stream emerging from the low mass, main sequence secondary star's Roche lobe is entrained onto the WD's magnetic field lines and funneled directly onto its magnetic pole(s) \\cite[see review in][ch.\\ 6]{warner95}. Other distinguishing characteristics of the AM Her stars include WD spin periods synchronized with their orbital periods and a high degree of linear and circular polarization (leading to the alternate name ``polars'' for this class of CV). Occasional reductions or interruptions to the accretion flow cause AM Her systems to drop to ``low states'' of reduced brightness. The origin of this accretion modulation is not fully known but may be related to, for example, solar-type magnetic activity (starspots) on the secondary star \\citep{hessman00}.  Modeling of the cyclotron lines in the optical spectrum of VV Pup (which were first noted by \\citealt{visv79} and \\citealt{wick79}) allowed \\citet{barrett85} to estimate a magnetic field strength of $B=31.5$ MG at the accreting pole.  The pole is located at an azimuth (i.e., the angle between the line of centers of the component stars and the projection of the magnetic axis onto the orbital plane of the CV) of $\\psi\\approx50^{\\circ}$ and a colatitude (i.e., the angle between the rotation and magnetic axes of the WD; see \\citealt{warner95}, his Figure 6.3, for definitions of angles in polar geometry) of $\\delta\\approx150^{\\circ}$. Combined with a system inclination of $i\\approx75^{\\circ}$ for VV Pup, this pole is visible for $\\approx$45\\% of the CV's orbit (\\citealt{cropper88} and references therein). \\citet{wick89} later used observations obtained during a very high accretion state to detect a $B=56$ MG field at the second magnetic pole, which forms an off-center dipole with the weaker first pole. Despite being weaker, the location of the first pole on the synchronously rotating WD -- on the side facing the secondary star -- makes it the preferred site for the accretion flow. However, accretion onto the second pole (which is located within $\\approx10^{\\circ}$ of the WD rotation axis, so is always visible; \\citealt{wick89}) sometimes also occurs, as evidenced by polarimetric observations made when the first pole is not visible (due to orbital motion) and differences in the total system brightness when the first pole is inactive ($V\\approx16$ when the second pole is active vs.\\ $V\\approx18$ when both poles are inactive; e.g., \\citealt{liebert79}). A blackbody fit to the optical spectrum of VV Pup during an extended, steady low state (both poles inactive) in 1977 suggested a temperature of $\\approx9000$ K for the WD \\citep{liebert78}.  VV Pup has been extensively observed at infrared \\citep{szkody83}, optical \\citep{warner72, imamura00}, ultraviolet \\citep{patterson84}, extreme-ultraviolet (EUV; \\citealt{vennes95}), and X-ray \\citep{patterson84, imamura00}  wavelengths. The EUV data indicated a possible oxygen overabundance and the presence of a hot accretion region.  To further study these properties of VV Pup, we obtained the first far-ultraviolet (FUV) observation of this magnetic CV, with the {\\em Far Ultraviolet Spectroscopic Explorer} ({\\em FUSE\\/}) satellite. ",
        "conclusions": "Analysis of our {\\em FUSE} observations of VV Pup has revealed many of the FUV characteristics of this magnetic CV. The shape of the FUV light curve is consistent with that seen at other wavelengths (optical, EUV, X-ray). Likewise, the characteristics of the phase-resolved FUV spectra during the bright portion of VV Pup's orbit (when the accreting magnetic pole is visible) agree with expectations from the geometric model for this system presented by \\citet{patterson84}. These include an increase in the continuum level at later orbital phases (as the magnetic pole and its associated accretion region rotate into more direct view) and changes in the shape and velocity shift of the prominent \\ion{O}{6} emission lines that we interpret in the context of the lines originating in the high velocity, hot, inner region of the accretion funnel located close to the WD surface. A hot temperature ($T\\gtrsim90,000$ K) is favored for a simple model of the FUV spectral energy distribution in VV Pup, in terms of both the fit to the {\\em FUSE} and archival {\\em IUE} continuum in the wavelength range 1000--3000 \\AA, and the size of the accretion region estimated from the normalization of the model. Although we expect {\\em a priori} that observations in the FUV should detect a cooler (and more extended) accretion region than did the {\\em EUVE} observations of \\citet{vennes95}, we cannot rule out a temperature as high as the 300,000 K determined for the EUV as an adequate fit to the {\\em FUSE} data."
    },
    "0206/astro-ph0206440_arXiv.txt": {
        "abstract": "We present deep $BVI$ observations of the core of M35 and a nearby comparison field obtained at the WIYN 3.5m telescope under excellent seeing conditions.  These observations probe to $V >$ 26, and display the lower main sequence in $BV$ and $VI$ CMDs down to $V$ = 23.3 and 24.6, respectively.  At these faint magnitudes the background Galactic field stars are far more numerous than the cluster stars, yet by using a smoothing technique and CMD density distribution subtraction we are able to recover the cluster fiducial main sequence and luminosity function to $V = 24.6$.  We find the location of the M35 main sequence in these CMDs to be consistent with earlier work on other open clusters, specifically NGC 188, NGC 2420, and NGC 2477.  We compare these open cluster fiducial sequences to stellar models by Baraffe \\etal (1998), Siess \\etal (2000), Girardi \\etal (2000), and Yi \\etal (2001) and find that the models are too blue in both $B-V$ and $V-I$ for stars less massive than $\\sim0.4$ M$_{\\odot}$.  At least part of the problem appears to be underestimated opacity in the bluer bandpasses, with the amount of missing opacity increasing toward the blue.  M35 contains stars to the limit of the extracted main sequence, at M $\\approx$ 0.10--0.15 M$_{\\odot}$, suggesting that M35 may harbor a large number of brown dwarfs, which should be easy targets for sensitive near-IR instrumentation on 8--10m telescopes.  We also identify a new candidate white dwarf in M35 at $V = 21.36 \\pm 0.01$.  Depending on which WD models are used in interpreting this cluster candidate, it is either a very high mass WD ($1.05 \\pm 0.05$ M$_{\\odot}$) somewhat older (0.19--0.26 Gyr, 3--$4\\sigma$) than our best isochrone age (150 Myr), or it is a modestly massive WD (0.67--0.78 M$_{\\odot}$) much too old (0.42--0.83 Gyr) to belong to the cluster. Follow-up spectroscopy is required to resolve this issue. ",
        "introduction": "NGC 2168 (M35) is a rich open cluster with an age similar to the Pleiades.  Since M35 is more populous and covers a smaller angular extent than the Pleiades, it offers excellent opportunities for studies of stellar evolution at $\\sim100$ Myr, even though M35 is further away and suffers greater background contamination.  Astrometric studies of M35 have a long history (Ebbighausen 1942; Meurers \\& Schwarz 1960; Lavdovskij 1961; Cudworth 1971; McNamara \\& Sekiguchi 1986a), and in fact continue to this day--M35 is {\\it the} astrometric calibrator for the HST Fine Guidance Sensors (McArthur \\etal 1997). Modern photometric studies of this cluster begin with Sung \\& Lee (1992) who obtained photoelectric $UBV$ photometry for 112 field plus cluster stars to $V = 14$, approximately the same limiting magnitude as the two more recent proper motion studies. Sung \\& Lee derived a true distance modulus of 9.3, an age of 85 Myr, and internal differential reddening of $0.26 \\leq$ E($B-V$) $\\leq 0.44$.  In a subsequent study, Sung \\& Bessel (1999) obtained $UBVI$ CCD photometry for stars brighter than $V = 20$ in a central $20^{\\arcmin}.5 \\times 20^{\\arcmin}.5$ cluster field.  From these data they derived $(V-M_{V})_{\\rm o} = 9.6 \\pm 0.1$, E$(B-V) = 0.255 \\pm 0.024$ (corresponding to $(V-M_{V}) = 10.39$ for $R_{V} = 3.1$), log age = 8.3 $\\pm$ 0.3 (200 Myr), [Fe/H] $\\approx -0.3$ (based on $U-B$ color excess), a present day mass function slope of $-2.1 \\pm 0.3$,\\footnote{All mass function slopes presented here are on the system $n(m) \\propto m^{-(1+x)}$, where the reported slope = $-(1+x)$, and the Salpeter (1955) value is $-2.35$.} and a binary frequency $\\geq 35$\\%.  A younger cluster age of 70 to 100 Myr was found by Reimers \\& Koester (1988a), based on a reanalysis of older photometry along with isochrones from Maeder \\& Mermilliod (1981) and on the cooling age of two cluster white dwarfs. Barrado y Navascu\\'es, Deliyannis, \\& Stauffer (2001a) have derived the cluster metallicity, [Fe/H] = $-0.21 \\pm 0.10$, from high resolution spectroscopy.  In the most recent photometric study of M35, Barrado y Navascu\\'es \\etal (2001b, hereafter BSBM), using the Kitt Peak 4m and CFH 3.6m telescopes, imaged the central $28{\\arcmin} \\times 28{\\arcmin}$ of the cluster in $VRI$ to $V \\approx 22$ and $I \\approx 23$.  BSBM found a luminosity function similar to the Pleiades, with a peak near $M_{I} = 9$, and present-day mass function characterized by three different power law slopes over the mass range 1.6 to $\\sim0.1$ M$_{\\odot}$.  BSBM found their central cluster field to contain $\\sim1600$ M$_{\\odot}$ among cluster members.  In a dynamical study of M35, Leonard \\& Merritt (1989) found that M35 is close to dynamical equilibrium, that its dynamical mass within the central 3.75 pc is 1600 to 3200 M$_{\\odot}$ (95\\% confidence), and that its IMF slope is $-2.7 \\pm 0.4$ between 1 and 6 M$_{\\odot}$.  Mathieu (1983), McNamara \\& Sekiguchi (1986b), and BSBM all noted that M35 exhibits mass segregation, though it is unclear whether this is due to relaxation or initial conditions.  Mathieu (1983) pointed out that the cluster age is close to the expected relaxation time of the intermediate mass component, though the relaxation time scale is uncertain by a factor of 2.  We obtained deep $BVI$ photometry of M35 in order to study the low mass main sequence stars and to search for cluster white dwarfs.  Our study presents higher signal-to-noise data for the faintest stars than the BSBM study, and we achieve this depth in $B, V$, and $I$, whereas BSBM achieve their greatest depth in $R$ and $I$.  Our photometry allows us to isolate the fiducial main sequence of the cluster in $B$ and $V$, as well as $V$ and $I$, which we compare to a range of stellar models.  The smaller field of view of our study precludes a detailed luminosity function or mass function study, however, as done by BSBM.  In these trade-offs between field of view and depth in various filters, our two deep photometric studies are complementary.  In addition, we have found a candidate cluster white dwarf which, if a {\\it bona fide} cluster member, places constraints on a combination of cluster age and stellar evolution. ",
        "conclusions": "We obtained deep $BVI$ observations of M35 and a nearby comparison field with the WIYN 3.5m telescope under non-photometric but excellent seeing conditions.  We calibrated the data against shallower 0.9m data (from Deliyannis \\etal 2002), achieving a photometric accuracy of approximately 0.02 mag.  These deep observations display the lower main sequence in the $BV$ and $VI$ CMDs down to $V$ = 23.3 and 24.6, respectively.  At these faint magnitudes the background Galactic field stars are far more numerous than the cluster stars, yet by using a smoothing technique (Silverman 1986) and CMD density distribution subtraction we were able to recover the cluster fiducial main sequence and luminosity function to $V = 24.6$.  We find the location of the M35 main sequence in these CMDs to be consistent with earlier work on other open clusters, specifically NGC 188 (WOCS1), NGC 2420 (von Hippel \\& Gilmore 2000), and NGC 2477 (von Hippel \\etal 1996).  On comparing these open cluster fiducial sequences to stellar models by Baraffe \\etal (1998), Siess \\etal (2000), Girardi \\etal (2000), and Yi \\etal (2001) we find that the models are too blue in both $B-V$ and $V-I$ for stars less massive than $\\sim0.4$ M$_{\\odot}$.  At least part of the problem appears to be underestimated opacity in the bluer bandpasses, with the amount of missing opacity increasing toward the blue.  M35 contains stars to the limit of the extracted main sequence, at M $\\approx$ 0.10--0.15 M$_{\\odot}$, suggesting that M35 may harbor a large number of brown dwarfs.  These brown dwarfs should be easy targets for sensitive near-IR instrumentation now being mounted on 8--10m telescopes. In fact, imaging observations of only one hour in $K$ would allow one to obtain S/N = 30 photometry 1 magnitude fainter than the brown dwarf limit in this cluster.  We also identify a new candidate white dwarf in M35 at $V = 21.36 \\pm 0.01$.  Depending on which WD models are used in interpreting this object, it is either a very high mass WD ($1.05 \\pm 0.05$ M$_{\\odot}$) somewhat (3--$4\\sigma$) older than our best isochrone age (150 Myr), or it is a modestly massive WD (0.67--0.78 M$_{\\odot}$) much too old (0.42--0.83 Gyr) to belong to the cluster.  Follow-up spectroscopy is required to resolve this issue."
    },
    "0206/astro-ph0206395_arXiv.txt": {
        "abstract": "Employing hydrodynamic simulations of structure formation in a $\\Lambda$CDM cosmology, we study the history of cosmic star formation from the ``dark ages'' at redshift $z \\sim 20$ to the present.  In addition to gravity and ordinary hydrodynamics, our model includes radiative heating and cooling of gas, star formation, supernova feedback, and galactic winds.  By making use of a comprehensive set of simulations on interlocking scales and epochs, we demonstrate numerical convergence of our results on all relevant halo mass scales, ranging from $10^8$ to $10^{15}\\, h^{-1}{\\rm M}_\\odot$.  The predicted density of cosmic star formation, $\\dot \\rho_\\star(z)$, is broadly consistent with measurements, given observational uncertainty.  From the present epoch, $\\dot \\rho_\\star(z)$ gradually rises by about a factor of ten to a peak at $z\\sim 5-6$, which is beyond the redshift range where it has been estimated observationally. In our model, fully 50\\% of the stars are predicted to have formed by redshift $z\\simeq 2.14$, and are thus older than 10.4 Gyr, while only 25\\% form at redshifts lower than $z\\simeq 1$.  The mean age of all stars at the present is about 9 Gyr.  Our model predicts a total stellar density at $z=0$ of $\\Omega_\\star= 0.004$, corresponding to about 10\\% of all baryons being locked up in long-lived stars, in agreement with recent determinations of the luminosity density of the Universe.  We determine the ``multiplicity function of cosmic star formation'' as a function of redshift; i.e.~the distribution of star formation with respect to halo mass.  At redshifts around $z\\simeq 10$, star formation occurs preferentially in halos of mass $10^{8}-10^{10}\\, h^{-1}{\\rm M}_\\odot$, while at lower redshifts, the dominant contribution to $\\dot \\rho_\\star(z)$ comes from progressively more massive halos.  Integrating over time, we find that about 50\\% of all stars formed in halos less massive than $10^{11.5}\\,h^{-1}{\\rm M}_\\odot$, with nearly equal contributions per logarithmic mass interval in the range $10^{10}-10^{13.5}\\, h^{-1}{\\rm M}_\\odot$, making up $\\sim 70\\%$ of the total.  We also briefly examine possible implications of our predicted star formation history for reionisation of hydrogen in the Universe. According to our model, the stellar contribution to the ionising background is expected to rise for redshifts $z>3$, at least up to redshift $z\\sim 5$, in accord with estimates from simultaneous measurements of the H and He opacities of the Lyman-$\\alpha$ forest. This suggests that the UV background will be dominated by stars for $z>4$, provided that there are not significantly more quasars at high-z than are presently known.  We measure the clumping factor of the gas from the simulations and estimate the growth of cosmic H{\\small II} regions, assuming a range of escape fractions for ionising photons.  We find that the star formation rate predicted by the simulations is sufficient to account for hydrogen reionisation by $z\\sim 6$, but only if a high escape fraction close to unity is assumed. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\footnotetext[1]{E-mail: volker@mpa-garching.mpg.de} \\footnotetext[3]{\\hspace{0.03cm}E-mail: lars@cfa.harvard.edu}  Hierarchical galaxy formation \\citep{Whi78} within a $\\Lambda$CDM cosmogony is currently the most successful paradigm for understanding the distribution of matter in the Universe.  In this scenario, structure grows via gravitational instability from small perturbations seeded in an early inflationary epoch.  The dominant mass component is (unidentified) collisionless cold dark matter, which also determines the dynamics of the baryons on large scales, where hydrodynamic forces are unimportant compared to gravity.  The successes of the $\\Lambda$CDM model are impressively numerous, ranging from a detailed picture of the primary anisotropies of the CMB at high redshift to the clustering properties of galaxies in the local Universe.  However, while the dynamics of collisionless dark matter is quite well understood, the same cannot be said for the baryonic processes that are ultimately responsible for lighting up the Universe with stars. Hydrodynamic simulations have been shown to produce reliable results for gas of low to moderate overdensity, allowing for detailed theoretical studies of e.g. the Lyman-alpha forest \\citep{CMOR94,Zh95,He96} and the intergalactic medium \\citep[e.g.][]{Dave2001,Cr00,Keshet2002}, but at gas densities sufficiently high for star formation to occur, our knowledge is much less certain.  Despite numerous attempts to include star formation in cosmological simulation \\citep[e.g.][]{Ce93,Ce00,Ka96,Kat99, Ye97,We97,Wein2000,St95,Bl99,Pea99,Pea2000,Th98,Th2000}, progress has been comparatively slow with this approach, primarily because of the physical complexity of star formation and feedback, but also because of the computational difficulties inherent in the task of embedding these processes into the framework of hierarchical galaxy formation.  For this reason, the most popular approach for describing galaxy evolution in hierarchical universes has been the so-called semi-analytic technique \\citep[e.g.][]{Wh91,Col91,Lac91,Lac93b,Kau93a,Kau94,Col94}. This method combines our firm knowledge of the dynamics and growth of dark matter halos with simplified parameterisations of the baryonic physics essential to galaxy formation.  In this manner, technical limitations of direct hydrodynamical simulations can be overcome, allowing some of the consequences of the assumed physics to be analysed.  On the other hand, the validity of the assumptions underlying the semi-analytic approach it is not always clear and must be confirmed by direct hydrodynamic simulations.  Moreover, while semi-analytic techniques are computationally much less expensive than hydrodynamical simulations, they are quite limited in their ability to relate galaxies to the surrounding IGM.  One strength of hydrodynamical simulations is that they can explore this relationship {\\em from first principles}, making it possible to constrain galaxy formation and evolution using the rich body of observational data on the IGM, as probed by quasar absorbers. Motivated by our desire to understand the connection between galaxies and their environments in detail, we here attempt to refine the methodology of direct simulations of galaxy formation, using a novel approach to describe star formation and feedback.  In this paper, we focus on one aspect of galaxy evolution; namely the global history of star formation in the Universe.  This is currently one of the most fundamental quantities in observational and theoretical cosmology, and it provides a crucial test for any theory of galaxy formation.  Over the last decade, observational results for the redshift evolution of the cosmic star formation rate density, $\\dot \\rho_\\star(z)$, have become available, at both low and high redshift \\citep[e.g.][]{Gal95,Mad96,Mad98,Lil96,Cow96,Cow99,Con97,Hug98,Trey98, Tresse98,Pas98,Steid99,Flo99,Gron99,Bald02,Lan2002}.  While these measurements are still fraught with large observational uncertainty, they are nevertheless beginning to constrain the epoch of galaxy formation.  In recent years, therefore, a number of studies have attempted to compute $\\dot \\rho_\\star(z)$ theoretically and to compare it with data, using either semi-analytic models \\citep{Bau98,Som00}, or numerical simulations \\citep{Pea2000,Nag00,Nag01,Asc02}. However, because of the complexity of the relevant physics and the large range in scales involved, the theoretical predictions have remained quite uncertain, a situation we try to improve on in this study.  Besides gravity, ordinary hydrodynamics, and collisionless dynamics of dark matter, our numerical simulations include radiative cooling and heating in the presence of a UV background radiation field, star formation, and associated feedback processes.  Cooling gas settles into the centres of dark matter potential wells, where it becomes dense enough for star formation to occur.  We regulate the dynamics of the gas in this dense interstellar medium (ISM) by an effective multi-phase model described by \\citet{SprHerMultiPhase}. In this manner, we are able to treat star formation and supernova feedback in a physically plausible and numerically well-controlled manner, although there are clearly uncertainties remaining with respect to the validity and accuracy of our description.  In principle, the cooling-rate of gas in halos can be computed accurately within simulations, and it is this rate that ultimately governs the global efficiency of star formation, provided that strong feedback processes are unimportant.  However, it is clear \\citep{Wh91} that cooling by itself is so efficient that it yields a collapse fraction of baryons that is considerably higher than that implied by the measured luminosity density of the Universe \\citep[e.g.][]{Bal01}. Simulations and semi-analytic models agree in this respect \\citep{Yoshida2001}.  Feedback mechanisms related to star formation are commonly invoked to reduce the efficiency of star formation.  In our model, we include such a strong feedback process in the form of galactic winds emanating from star-forming regions. There is observational evidence for the ubiquitous presence of such winds in star forming galaxies both locally and at high redshift \\citep[e.g.][]{Heck95,Bland95, Lehn96,Dahlem97,Mar99,Heck00,Frye01,Pett00,Pett01}.  It is believed that these galactic outflows are powered by feedback energy from supernovae and stellar winds in the ISM, but their detailed formation mechanism is not entirely clear.  Since we also lack the ability to spatially resolve the interactions of supernova blast waves and stellar winds within the ISM, we invoke a phenomenological model for the generation of galactic ``superwinds''.  Note, however, that the hydrodynamic interaction of the wind with infalling gas in the halos and with the IGM is treated correctly by our numerical code.  In particular, we are thus able to investigate how winds associated with star formation influence galaxy formation, how they disperse and transport metals, and how they heat the intergalactic medium.  We base our work on a large set of numerical simulations that cover a vast range of mass and spatial scales.  For example, the masses of resolved halos in our study spans more than a factor of $10^9$. Our simulation programme was designed to examine star formation on essentially all relevant cosmological scales, enabling us to arrive at a reliable prediction for the star formation density over its full history, ranging from the present epoch far into the ``dark ages'' at very high redshift.  As an integral part of our simulation set, we also carried out extensive convergence tests, allowing us to cleanly quantify the reliability of our results.  As part of our analysis, we introduce a ``multiplicity function of star formation'' which gives the cumulative star formation density per logarithmic interval of halo mass at a given epoch.  Using this quantity, the global star formation density can be decomposed into the number density of halos of a given mass scale (i.e. the halo ``mass function''), and the average star formation efficiency of halos of a given mass.  Because the cosmological mass function has been reliably determined using large collisionless simulations \\citep{Jen01}, this decomposition allows us to nearly eliminate the dependence of our results on cosmic variance.  We also briefly investigate the potential relevance of our results for the reionisation of hydrogen in the Universe. It is a long-standing question which sources dominate the ionising UV flux as a function of redshift.  While the ratio of He{\\small II} and H{\\small I} optical depths observed at $z\\sim 2.4$ indicates that quasars are the most significant source of ionising radiation at this low redshift, it has been suggested that massive stars might dominate at higher redshift \\citep[e.g.][]{Haehn2001,Sok2002b}.  Using our results, we test the possibility that high-$z$ star formation by itself could have been responsible for reionisation of hydrogen at a redshift of around $z\\sim 6$.  This paper is organised as follows. In Section~\\ref{SecMethod}, we describe our simulations, and the analysis applied to them.  We then move on to discuss tests for numerical convergence in Section~\\ref{SecConv}.  In Section~\\ref{SecMulti}, we introduce the concept of a multiplicity function for cosmic star formation, which we use in Section~\\ref{SecSFR} to derive our composite result for the cosmic star formation history, and compare it to observational constraints.  In Section~\\ref{SecWhenWhere}, we analyse consequences of our prediction for the age distribution of stars and then discuss possible implications for the reionisation of the Universe in Section~\\ref{SecReion}.  Finally, we summarise our findings in Section~\\ref{SecDisc}. ",
        "conclusions": "\\label{SecDisc}  We have studied the star formation history of the Universe using cosmological SPH simulations that employ a sophisticated treatment of star formation, supernova feedback, and galactic outflows.  Using present day computational capabilities, we have been able to show that this model yields numerically converged results for halos of all relevant mass scales, assuming that the vast majority of stars can form only in halos in which gas can condense by atomic line cooling. Using a large set of high-resolution simulations on interlocking scales and at interlocking redshifts, we have been able to determine the multiplicity function of cosmic star formation and infer its evolution from high redshift to the present.  We think that the numerical methodology used here represents perhaps the most comprehensive attempt thus far to constrain the cosmic star formation history in the $\\Lambda$CDM cosmogony with simulations.  Interestingly, our predicted star formation density peaks early, already at redshifts around $z\\sim 5 - 6$, and then falls to the present time by a factor of about 10. The decline at low redshift is hence shallower than suggested by a subset of observational results, and our star formation rate also appears low at redshifts around $z\\sim 1$ compared to some of the data. However, there is very large scatter between different observations, which are partly inconsistent with each other, emphasising the challenges posed by these measurements.  Our result matches a number of the observational points, and does so particularly well at very low and very high redshift.  We thus conclude that at present our results are consistent with direct data on the star formation rate density as a function of epoch.  It will be very interesting to see whether this will remain true when observational uncertainty is reduced in the future.  One consequence of the star formation history we predict is that the majority of stars at the present should be quite old, despite the hierarchical formation of galaxies. Already 10.4 Gyr ago, by redshift $z=2.14$, half of the stars should have formed, with only $\\sim 25\\%$ forming at redshifts below unity. The mean stellar age at $z=0$ predicted by our model is 9 Gyr.  Integrated over cosmic time, star formation occurs predominantly in halos with masses between $10^{8}\\,h^{-1}{\\rm M}_\\odot$ and $10^{14}\\,h^{-1}{\\rm M}_\\odot$, with 50\\% forming in halos below $10^{11.5}\\,h^{-1}{\\rm M}_\\odot$.  It is thus clear that simulations of galaxy formation need to be able to resolve at least these mass scales well in order to have a chance at giving a reasonably accurate accounting of the formation of the luminous component of the Universe.  The total integrated star formation rate predicts a density in stars of about $\\Omega_\\star=0.004$, or expressed differently, 10\\% of all baryons should have been turned into long-lived stars by the present. This is in comfortable agreement with recent determinations of the luminosity density of the Universe, while earlier theoretical work was typically predicting substantially larger numbers of stars, by up to a factor of three or so.  Our model hence appears to have resolved the so-called ``over-cooling crisis''.  This was primarily made possible by the strong feedback we adopted in our simulations in the form of galactic winds.  We note that our predictions do depend on the model for star formation and feedback we adopted.  In particular, without the inclusion of galactic outflows, which have been introduced on a phenomenological basis in our approach, star formation in the low redshift universe would clearly have been higher. Our strategy has been to normalise the free parameters in our star formation law (the consumption timescale of cold gas) to observations of local disk galaxies, and to select the parameters for the galactic winds as suggested by observations.  Under the assumption that the same laws hold roughly at all redshifts, we have then computed what simulations predict for the $\\Lambda$CDM model.  In this context, one should clearly distinguish between the computational difficulty of the problem on one hand, and the uncertainty and complexity of the modeling of the physics on the other.  We think that we have been able to make great progress on the computational side of the problem, but we are aware that large uncertainties remain in our handling of star formation and feedback. It is possible that the modeling of the physics we adopted could be incomplete in crucial respects.  Perhaps one of the most important effects that has been neglected in our simulations is metal line cooling.  It is well known that metals can substantially boost cooling rates, and hence can potentially have a very prominent effect on the rates at which gas becomes available for star formation \\citep{Wh91}.  However, the extent to which metal enrichment can enhance gas cooling is strongly dependent on how efficiently metals can be dispersed and mixed into gas that has yet to cool for the first time.  In simulations without galactic winds, we find that metals are largely confined to the star-forming ISM at high overdensity.  Metal-line cooling would not significantly increase the star formation rates in these simulations.  In the present set of simulations, metals can be transported by winds into the low-density IGM.  If the corresponding gas is reaccreted at later times into larger systems, metal-line cooling should then accelerate cooling, thereby potentially increasing the total star formation density compared to what we estimated here, particularly at low redshift.  But recall that metal line cooling has also been neglected when we normalised the star formation timescale used in our model for the ISM to match the Kennicutt law. If we had included metal line cooling, this normalisation would have been somewhat different in order to compensate for the accelerated cooling, such that the net star formation rate would have again matched the Kennicutt law. This effect will alleviate any difference that one naively expects from the inclusion of metal line cooling.  More work is therefore needed to quantitatively estimate how important the effect of metal line cooling would ultimately be in the present model of feedback due to galactic winds.  The set of simulations we carried out offers rich information on many aspects of galaxy formation and structure formation, not only on the cosmic star formation history.  Note in particular that our simulations are among the first that can self-consistently address the interaction of winds with the low-density IGM, and the transport of heavy metals along with them.  We have already investigated effects of the winds on secondary anisotropies of the cosmic microwave background \\citep{WhiHerSpr02}.  Among the issues we plan to address next, is the question of whether winds imprint specific signatures in the Lyman-$\\alpha$ forest that can be identified observationally.."
    },
    "0206/astro-ph0206506_arXiv.txt": {
        "abstract": "s{ The INTEGRAL Science Data Centre (ISDC) processes, archives and distributes data from the INTEGRAL mission. At the ISDC incoming data from the satellite are processed and searched for transient sources and Gamma-Ray bursts. The data are archived and distributed to the guest observers. As soon as the data are public, any astronomer can access the data via the internet. ISDC also provides the tools which are necessary for the data analysis and offers user support concerning questions related to the INTEGRAL data. ISDC acts as a contacting point between the scientific community and the various instrument teams. In this proceeding an example for SPI data processing is shown.} ",
        "introduction": "Data from the INTEGRAL mission (see e.g. Winkler \\& Hermsen~\\cite{integral}), which is going to be launched October 2002, will be made available to the scientific community via the INTEGRAL Science Data Centre (ISDC; Courvoisier {\\em et al.}~\\cite{ISDC}). The ISDC is hosted by the Geneva observatory and is funded by an international consortium with ESA support. The ISDC is the contact point between the community, instrument teams, INTEGRAL Science Operations Centre (ISOC), Mission Operations Centre (MOC), the INTEGRAL project and the Russian data center. In this contribution we will discuss briefly the services provided by the ISDC, in terms of data and software which is made available, and show an example how the user can perform a scientific analysis using the INTEGRAL data with ISDC software. ",
        "conclusions": ""
    },
    "0206/hep-ph0206211_arXiv.txt": {
        "abstract": "\\vspace{1cm} We investigate whether present data on helioseismology and solar neutrino fluxes may constrain WIMP--matter interactions in the range of WIMP parameters under current exploration in WIMP searches. We find that, for a WIMP mass of 30 GeV, once the effect of the presence of WIMPs in the Sun's interior is maximized, the squared isothermal sound speed is modified, with respect to the standard solar model, by at most 0.4\\% at the Sun's center. The maximal effect on the $^8$B solar neutrino flux is a reduction of 4.5\\%. Larger masses lead to smaller effects.  These results imply that present sensitivities in the measurements of solar properties, though greatly improved in recent years, do not provide information or constraints on WIMP properties of relevance for dark matter. Furthermore, we show that, when current bounds from direct WIMP searches are taken into account, the effect induced by WIMPs with dominant coherent interactions are drastically reduced as compared to the values quoted above. The case of neutralinos in the minimal supersymmetric standard model is also discussed. ",
        "introduction": "\\label{sec:intro}  A host of independent astronomical observations point to the existence in our Universe of a total amount of matter in the range $0.2 \\lsim \\Omega_m \\lsim 0.4$ \\cite{matter} (or equivalently, $0.05 \\lsim \\Omega_m h^2 \\lsim 0.3$, where $h$ is the Hubble constant in units of 100 km $\\cdot$ s$^{-1} \\cdot$ Mpc$^{-1}$), well beyond the amount of visible matter $\\Omega_{vis} \\sim 0.003$.  Since the primordial nucleosynthesis tells us that baryons cannot contribute for more than about 5\\% \\cite{bbn}, most of the dark matter must be non-baryonic. The evolution theory of the primordial density fluctuations into the present cosmological structures indicates that most of the dark matter must be comprised of cold particles, {\\it i.e.} of particles that decoupled from the primordial plasma when nonrelativistic. A particle with the suitable properties for being a significant cold dark matter relic is generically defined as a WIMP (weakly interacting massive particle).  A variety of  physical realizations for a WIMP are offered by various extensions of the Standard Model  \\cite{nic}, the neutralino being one of the most appealing candidates. In the present paper, many considerations are developed in terms of generic WIMPs. We consider the neutralino, whenever we wish to narrow down to a specific candidate.  WIMPs are very actively searched for by means of various experimental strategies. Direct searches rely on the measurements of the signal that a nucleus of an  appropriate detector  would generate, when hit by a WIMP \\cite{morales}. Indirect  searches are based on measurements of the signals due to WIMP pair annihilations in the galactic halo or inside celestial bodies \\cite{nic}.  The event rates of WIMP direct search experiments are proportional to the product of the local WIMP density $\\rho_{\\chi}$ in the Galaxy times the WIMP-nucleus cross--section. In what follows $\\rho_{\\chi}$ will be expressed as $\\rho_{\\chi} = \\xi \\cdot \\rho_l$, where $\\rho_l$ is the local {\\it total} dark matter density and $\\xi$ ($\\xi \\leq 1$) is a scaling parameter which accounts for the actual fraction of local dark matter to be ascribed to the candidate $\\chi$.  WIMP--matter interactions are conveniently classified in terms of coherent cross--sections and (nuclear) spin--dependent ones.  In the first case, the WIMP--nucleus cross--section is given in terms of the WIMP--proton and WIMP--neutron cross--sections. To simplify the formalism, in the following we further assume that, in the coherent case, WIMPs interact with equal strength with protons and neutrons (for instance, this may safely be assumed for neutralinos, while it is not the case for neutrinos), and then we express a generic WIMP--nucleus coherent cross--sections in terms of a single WIMP--nucleon cross--section $\\sigma_c$.  For the spin--dependent case, the derivation of a WIMP--nucleon cross--section from the WIMP--nuclear one depends on the nature of the WIMP and on specific nuclear properties \\cite{bdmsbi}.  In the case of WIMPs with coherent interactions with matter, the range of WIMP parameters under current exploration in WIMP direct searches is conveniently expressed in terms of the quantity $\\xi \\sigma_c$: \\begin{equation} 10^{-43} \\; {\\rm cm^2} \\leq \\ \\xi \\sigma_c \\leq 6 \\cdot 10^{-41} \\; {\\rm cm^2} \\, , \\label{eq:sens} \\end{equation} for WIMP masses $m_{\\chi}$ in the range: \\begin{equation} 30 \\; {\\rm GeV} \\leq  m_\\chi \\leq 270 \\;  {\\rm GeV}. \\label{eq:mass} \\end{equation}  In the derivation of the ranges in Eqs. (\\ref{eq:sens}--\\ref{eq:mass}), one has taken into account a variety of WIMP distribution functions in the galactic halo and uncertainties in the determination of the relevant astrophysical quantities \\cite{belli1,belli2}.  For the reasons mentioned above, in the case of spin--dependent interactions no model--independent sensitivity range may be derived. However, to give an indication, we recall that, in the case of a relic particle which interacts with matter mainly by spin-dependent interactions mediated by $Z$--boson exchange, current WIMP direct experiments are sensitive to values of $\\xi \\sigma_{sd}$ of about $10^{-36} - 10^{-37}$ cm$^2$ \\cite{bdmsbi}, where $\\sigma_{sd}$ is the spin--dependent WIMP-proton cross--section.  One of the WIMP direct search experiments, the DAMA/NaI(Tl) experiment \\cite{dama}, has observed an annual modulation effect (at a $4\\sigma$ C.L.) with all the features expected for a WIMP signal \\cite{freese}. When interpreted as due to a WIMP with coherent interactions, the DAMA effect provides a $3\\sigma$ region in the plane $m_{\\chi} - \\xi \\sigma_c$ embedded in the range of Eqs.(\\ref{eq:sens}--\\ref{eq:mass}).  In Ref. \\cite{noi} it was proved that this annual modulation effect is compatible with an interpretation in terms of relic neutralinos (see also Ref.\\cite{others}).  If WIMPs populate our Galaxy, they can be captured by the Sun and accumulate in its core. For suitable values of the WIMP parameters, this could have relevant effects on the Sun. In fact, WIMPs could provide an effective mechanism for energy transport in the Sun, producing an isothermal core and reducing substantially the Sun central temperature $T_c$.  Following this idea, some time ago a special class of WIMPs named {\\it cosmions}, with masses of a few GeV and scattering cross sections on nucleons of the order of $10^{-36}$ -- $10^{-34}$ cm$^2$, was studied in detail as a way of solving simultaneously the solar neutrino puzzle and the dark matter problem (see e.g. Refs. \\cite{Gould:hm,Spergel:1985,Gould:1989ez,Gilliland:1986,Dearborn:1990mm,cosmions_others}). The cosmion hypothesis was progressively abandoned when it became clear that the solar neutrino puzzle cannot be accounted for by simply reducing $T_{c}$.  In the last twenty years, our observational knowledge of the solar interior has progressed enormously. By means of helioseismic data it has become possible to derive the sound speed with an accuracy of about one part per--thousand over most of the radial profile and of about one percent in the innermost part \\cite{Ricci:1997}. By the same method, it has been possible to deduce important properties of the convective envelope.  The photospheric helium fraction $Y_{ph}$ and the depth of the convective envelope $R_{b}$ have been determined with an accuracy of about one per cent and one per--thousand respectively, following the pioneering papers of Refs. \\cite{Dziembowski:1991,Christensen:1991}.  Moreover, the measurement of the neutrino flux from $^8$B decay obtained by combining \\cite{Villante:1998pe,Fiorentini:2001jt,Fogli:2001vr,Fogli:2002ts} SNO charged current \\cite{Ahmad:2001an} and Super-Kamiokande \\cite{Fukuda:2001nj} data and confirmed by the recent SNO neutral current results \\cite{Ahmad:2002jz}, has provided a determination of the temperature $T_{c}$ near the center of the Sun at the level of about one per cent \\cite{Fiorentini:2001et}.  All the predictions of the Standard Solar Model (SSM) have been confirmed by these accurate tests, so one is naturally led to the question of whether our accurate knowledge of the Sun interior can be used to constrain the WIMP parameter space.  This particularly interesting question was recently raised in Refs. \\cite{Lopes:2001ra} and \\cite{Lopes:2001ig}, where it was concluded that solar physics can be used to significantly constrain the WIMP parameter space.   In this paper, we provide an alternative analysis of the problem, obtaining results substantially different from those derived in Refs. \\cite{Lopes:2001ra,Lopes:2001ig}.  The plan of the paper is as follows. In Section II we discuss WIMP energy transport in the Sun, introducing the relevant physical parameters to be discussed in the subsequent sections. In Section III and IV WIMP energy transport is applied to the central solar structure, in order to determine the regions of the WIMP parameter space where the WIMP energy transport is most efficient. This leads to an optimal choice of WIMP parameters which is then used in Section V to calculate the solar sound speed profile and discuss the possible impact of helioseismology on WIMP properties. The sensitivity of the available data on $^8$B solar neutrinos to the presence of WIMPs in the Sun is then discussed in Section VI. The neutralino, as a specific realization of WIMP, is discussed in Section VII. Section VIII is devoted to our conclusions ",
        "conclusions": "\\label{sec:finale}  Our understanding of solar physics has significantly advanced in recent times, because of remarkable improvements in helioseismology and of new data on solar neutrinos. In the present paper we have addressed the question of whether these new developments put solar physics in a position to provide  constraints on the possible presence of WIMPs in the core of the Sun, for values of WIMP parameters under current exploration in WIMP searches. We summarize here our main results: \\begin{itemize} \\item We have provided a quantitative criterium to determine whether a putative WIMP candidate could produce observable modifications of the solar structure. Namely, we have introduced the parameter $\\delta$, defined as the ratio between the WIMP luminosity and the radiative luminosity at the center of the Sun (see Eq.(\\ref{eq:delta})), which can be analytically calculated as a function of the WIMP-nucleon scattering cross section $\\sigma_p$ and the number of WIMPs in the Sun $N_{\\chi}$.  By considering the uncertainty in the radiative opacity, one finds that $\\delta \\gsim 10^{-2}$ is a necessary (but not sufficient) condition for WIMPs to have observable effects on the Sun. \\item We have calculated the number of WIMPs in the Sun for a generic WIMP candidate as a function of the WIMP-nucleon scattering cross section $\\sigma_p$ and of the WIMP-antiWIMP pair annihilation cross-section $\\sigma_a$, both for coherent and spin-dependent WIMP interactions. We have shown that a value $\\delta \\sim 10$ can be reached, if WIMP-antiWIMP annihilation is negligible and $\\sigma_p \\sim 10^{-37} \\;{\\rm cm}^2$ (coherent scattering) or $\\sigma_p \\sim 10^{-34} \\;{\\rm cm}^2$ (spin-dependent interactions).  For coherent WIMP scattering, such large values for the cross sections are excluded by direct search experiments, unless WIMPs give a subdominant contribution to dark matter of our galaxy.  In this case, if one takes into account direct search bounds ($\\xi\\sigma_p\\lsim 10^{-41} \\;{\\rm cm}^2 $, where $\\xi$ is the rescaling parameter which accounts for the actual fraction of local dark matter to be ascribed to the given WIMP), one obtains $\\delta\\lsim 10^{-6}$ in the assumption of no rescaling ($\\xi =1$) and $\\delta\\lsim 10^{-3}$ in the general case ($\\xi \\leq 1$). \\item We have considered the neutralino as a specific WIMP candidate. We have calculated the number of neutralinos in the Sun for MSSM configurations compatible with present experimental and cosmological constraints.  For each configuration, we determined the rescaling parameter $\\xi$ by the standard rescaling procedure described in Eq.~(\\ref{eq:rescal}).  In the case of no rescaling ($\\xi=1$), we obtained $\\delta\\lsim 10^{-9}$. For subdominant neutralinos ($\\xi\\leq 1$), we obtained $\\delta\\lsim 10^{-6}$ for configurations in which the capture rate is dominated by coherent scattering and $\\delta\\lsim 10^{-4}$ for configurations in which the dominant contribution is due to spin dependent interactions. \\item The previous points already show that solar physics is not competitive with direct experiments in a large part of the WIMP parameter space under current exploration in WIMP searches.  In order to complete our analysis and to understand whether some information could be obtained from the present helioseismic and solar neutrino data, we have constructed solar models, for WIMP masses above 30 GeV, choosing the value of the WIMP parameters which maximize the effect of WIMPs on the Sun ($\\delta \\sim 10$).  As a result of the presence of WIMPs in the Sun, we obtained variations of the sound speed profile and of the $^8$B neutrino flux which are within the current experimental uncertainties.  The smallness of the effects is essentially due to the the smallness of the WIMP estension region which, for WIMPs with masses larger than tens of GeV, is $r_\\chi \\leq 0.02 R_{\\odot}$. \\end{itemize}  In conclusion, no constraints can be derived at present from solar physics for WIMPs with masses above 30 GeV. Our conclusions are at variance with results derived in Refs. (\\cite{Lopes:2001ra,Lopes:2001ig}). The origins of these disagreements have been elucidated in the present paper.    \\appendix"
    },
    "0206/astro-ph0206038_arXiv.txt": {
        "abstract": "We suggest that molecular cloud (MC) turbulence is a consequence of the very process of MC formation by collisions of larger-scale flows in the diffuse atomic gas, which generate turbulence in the accumulated gas through bending-mode instabilities. Turbulence is thus maintained for as long as the accumulation process lasts ($\\sim$ several Myr). Assuming that supersonic turbulence in MCs has the double role of preventing global collapse while promoting the formation of smaller-scale structures by turbulent compression (``turbulent fragmentation''), we then note the following properties: a) Turbulent fragmentation necessarily deposits progressively smaller fractions of the total mass in regions of progressively smaller sizes, because the smaller structures are subsets of the larger ones. b) The turbulent spectrum implies that smaller scales have smaller velocity differences. Therefore, below some scale, denoted $l_{\\rm eq}$, the turbulent motions become subsonic. This is an energy {\\it distribution} phenomenon, not a dissipative one. On this basis, we propose that the star formation efficiency (SFE) is determined by the fraction of the total mass that is deposited in clumps with masses larger than $M_{\\rm J}(l_{\\rm eq})$, the Jeans mass at scale $l_{\\rm eq}$, because subsonic turbulence cannot promote any further subfragmentation. In this scenario, the SFE should be a monotonically increasing function of the sonic and turbulent equality scale, $l_{\\rm eq}$. We present preliminary numerical tests supporting this prediction, and thus the suggestion that (one of the) relevant parameter(s) is $l_{\\rm eq}$, and compare with previous proposals that the relevant parameter is the energy injection scale. ",
        "introduction": "Two of the main questions concerning molecular cloud (MC) structure and star formation are a) what is the origin and supply of MC turbulence? and b) what is the origin of the low efficiency of star formation? Indeed, MCs are known to be turbulent, with motions that are supersonic at scales $\\ga 0.1$ pc (Zuckerman \\& Evans 1974; Larson 1981). However, recent numerical simulations (Padoan \\& Nordlund 1999; Mac Low et al.\\ 1998; Stone,Ostriker \\& Gammie 1998; Mac Low 1999; see also Avila-Reese \\& V\\'azquez-Semadeni 2001 for the case of the global ISM) have suggested that strong MHD turbulence decays rapidly even in the presence of strong magnetic fields. Thus, it appears that MC turbulence must be continually driven during the entire lifetime of the clouds. The driving mechanism, however, is probably not restricted to the stellar activity {\\it internal} to the clouds, since clouds devoid of stars exhibit similar turbulent properties as clouds with stars (e.g., McKee 1999 and references therein).  Concerning the star formation efficiency (SFE), it is well known that it is low, on the order of a few percent with respect to the total cloud mass. Traditional explanations for this low efficiency have been, in the case of low-mass, isolated star formation, that the MC cores in which stars form are magnetically supported (i.e., ``subcritical''), so that collapse is delayed until the magnetic field diffuses out of the core by ambipolar diffusion (see, e.g., Shu, Adams \\& Lizano 1987) However, there exist several recent suggestions that all cores are critical or supercritical (e.g., Nakano 1998; Hartmann, Ballesteros-Paredes \\& Bergin 2001; Bourke et al.\\ 2001; Crutcher, Heiles \\& Troland 2002). On the other hand, it is known that turbulence can prevent global gravitational collapse of a cloud when the energy injection scale is smaller than the Jeans length (L\\'eorat, Passot \\& Pouquet 1990; Klessen, Heitsch \\& Mac Low 2000, hereafter KHM00), while promoting ``fragmentation'', i.e., the formation of smaller-scale density substructures that can possibly undergo local collapse (Sasao 1973; Tohline, Bodenheimer \\& Christodolou 1987; Elmegreen 1993; V\\'azquez-Semadeni, Passot \\& Pouquet 1996; Padoan 1995; Padoan et al. 2001; KHM00). KHM00 have shown that the efficiency in numerical simulations (measured as the fraction of mass in collapsed objects) is large when the energy injection scale is larger than the Jeans length, and low otherwise. However, it is possible, as we suggest in \\S 2 that MC turbulence is part of a cascade coming from larger scales. Thus, the description in terms of an energy injection scale smaller than the MC itself is probably not optimal for real MCs. Also, this shows that the origin of MC turbulence and of the SFE are intimately related.  In the present paper, we propose that {\\it MC turbulence originates from the same process that forms the clouds} (\\S 2), and sets an upper limit to the SFE through the fraction of the total mass it deposits in regions with sizes such that the turbulent velocity dispersion becomes subsonic (so that no further subfragmentation can occur within them), and masses larger than their Jeans mass (\\S 3). ",
        "conclusions": "\\label{sec:discussion}  We have shown numerical evidence that $l_{\\rm eq}$ is a relevant parameter in setting an upper limit to the SFE (other processes, such as stellar energy feedback, may contribute as well). However, it is most likely that other parameters are also involved. In particular, as the spectral energy distribution of the turbulence varies, the sizes of the density structures are likely to change as well. Padoan (1995) made an attempt in this direction based on the probability distribution function (PDF) of the density field. However, the PDF contains no information of the mass spatial distribution. What is necessary is the density PDF {\\it parameterized by region size}. This is precisely the kind of information contained in the multifractal spectrum, and which we are in the process of incorporating into the treatment."
    },
    "0206/astro-ph0206332_arXiv.txt": {
        "abstract": "The Optical Monitor telescope (\\cite{ablustin-C2_27:mas01}) on XMM-Newton provides an exciting multi-wavelength dimension to observations of Active Galactic Nuclei. Here we present ultraviolet images, taken with the OM UVW2 filter (140-270 nm), of various Seyfert galaxies, some of which have never been observed in this waveband before. The images show UV emission from both the active nucleus and the host galaxy. The distribution of UV emission in the galaxy shows where star formation is occurring, thus giving us clues as to the evolution of the host galaxy and perhaps its relationship to the Seyfert Nucleus. ",
        "introduction": "\\label{ablustin-C2_27_sec:7314_7313}  The UV image of \\object{NGC~7314} (Figure~\\ref{ablustin-C2_27_7314_7313:fig1}; top left), a Seyfert 1.9 galaxy at z=0.00474, shows that its active nucleus is far less bright in UV than the spiral arms, where the UV emission mostly coincides with the locations of HII regions (and optical emission) in the galaxy. The Seyfert nucleus sits in a mostly UV-dark region at the centre of the galaxy. Overall, the spiral arms are observed to extend about 1\\arcmin \\, horizontally (as seen here) and 2.7\\arcmin \\, vertically. We also see in this image the spiral galaxy \\object{NGC~7313} (z=0.01915), bottom right, 4.3\\arcmin \\, away from the nucleus of \\object{NGC~7314}. The bright object at the top right is a Galactic star.  \\begin{figure*} \\begin{center} \\epsfig{file=ablustin-C2_27_fig1.eps, width=14cm} \\end{center} \\caption{UV image of \\object{NGC~7314} (top left) and \\object{NGC~7313} (bottom right); image rotated so that North is upwards and East is to the left (all the images in this paper are oriented the same way), image size approx. 6\\arcmin x 5\\arcmin. The angular resolution of the image is 1\\arcsec, except for the central 2\\arcmin \\, of \\object{NGC~7314} which is at 0.5\\arcsec \\, resolution } \\label{ablustin-C2_27_7314_7313:fig1} \\end{figure*} ",
        "conclusions": ""
    },
    "0206/astro-ph0206042_arXiv.txt": {
        "abstract": "We present new imaging measurements of 27 individual globular clusters in the halo of the nearby elliptical galaxy NGC 5128, obtained with the {\\it Hubble Space Telescope} STIS and WFPC2 cameras. We use the cluster light profiles to determine their structural parameters (core and half-light radii, central concentration, and ellipticity). Combining these with similar data for selected inner-halo clusters from Holland et al.~1999 (AAp, 348, 418), we now have a total sample of 43 NGC 5128 globular clusters with measured structural properties. We find that classic King-model profiles match the clusters extremely well, and that their various structural parameters (core- and half-light radius, central surface brightness, central concentration) fall in very much the same range as do the clusters in the Milky Way and M31. We find half a dozen bright clusters which show tentative evidence for ``extra-tidal light'' that extends beyond the nominal tidal radius, similar in nature to several such objects previously found in the Milky Way and M31; these may represent clusters being tidally stripped, or possibly ones in which anisotropic velocity distributions are important. We also confirm previous indications that NGC 5128 contains relatively more clusters with large ($\\epsilon > 0.2$) ellipticity than does the Milky Way.  Instead, the $\\epsilon-$distribution of the NGC 5128 clusters strongly resembles that of the old clusters in the LMC and also in M31. Finally, calculations of the cluster binding energies $E_b$ as defined by McLaughlin 2000 (ApJ, 539,618) show that the NGC 5128 clusters occupy the same extremely narrow region of the parametric ``fundamental plane'' as do their Milky Way counterparts. Our data are thus strongly consistent with the claim that the globular clusters in both NGC 5128 and the Milky Way are fundamentally the same type of object:  old star clusters with similar mass-to-light ratios and King-model structures. ",
        "introduction": "Globular star clusters have remarkably simple structures that are well approximated by isotropic, single-mass \\citet{kin66} models. In the multi-dimensional space of all their structural quantities such as scale radii, central concentration, surface brightness, velocity dispersion, mass-to-light ratio, and so forth, it is striking that real globular clusters in the Milky Way inhabit only a narrow region referred to as the fundamental plane \\citep[FP; see][]{djo95}.  Recently, \\citet{mcl00} has shown that a particularly simple way of expressing the FP is to note that any King model is completely specified by four input parameters such as total cluster luminosity $L$, central concentration $c =$ log $(r_t/r_c)$, mass-to-light ratio, and binding energy $E_b$.  Adding in the two strong empirical constraints that $M/L \\simeq$ const and $E_b \\sim L^2$ then requires the clusters to lie on a two-dimensional slice of this 4-space, leaving only two quantities ($c$ and $L$) to determine the residual scatter on this FP. In turn, the concentration $c$ is correlated with $L$, leaving the remarkable result that the structures of these clusters are fixed largely by just one major {\\sl internal} parameter, their total mass (or luminosity, at a given age).  The {\\sl external} environment, i.e.~its location in the Galactic tidal field, also has some influence on the FP parameters.  However, clusters may have formed under drastically different environmental conditions in other galactic environments (giant and dwarf ellipticals, starburst systems, galactic bulges and rings, etc.).  The structures of clusters in these other types of galaxies remain to be investigated. The only other large galaxy for which this kind of study has been carried out in detail is M31, a large disk galaxy much like the Milky Way, and perhaps not surprisingly, its globular clusters strongly resemble those of the Milky Way \\citep{fus94,hol97,bar02}.  Beyond the Local Group, the nearest large galaxy containing many globular clusters available for detailed measurement is NGC 5128, the giant elliptical at the center of the Centaurus group at a distance $d \\sim 4$ Mpc. More importantly, it is a very different type of galaxy than any in the Local Group, and quantitative study of its clusters holds considerable promise for adding constraints to formation modelling.  In addition, a point of special interest from the viewpoint of globular cluster structural studies is that, because of sheer population size, NGC 5128 has many clusters at the upper end of the globular cluster mass distribution ($\\gtsim 10^6 M_{\\odot}$); with a total population of perhaps 1900 clusters \\citep{har84}, it should have 40 to 50 clusters with $M_V \\ltsim -10$. By contrast, all the Local Group galaxies combined contain perhaps $\\sim 600$ globular clusters and thus have only a few with $M_V \\ltsim -10$. NGC 5128 thus gives us the opportunity, in a single galaxy, to explore the empirical FP relations at extremely high cluster mass approaching $10^7 M_{\\odot}$ ($M_V \\sim -12$).  Clusters in galaxies as distant as Centaurus appear barely nonstellar under typical ground-based imaging resolutions of $1''$ but can be much more well resolved with the cameras on HST.  The first clear demonstration that accurate structural profiles and King-model parameters could be obtained for these objects was provided by \\citet{har98}, who studied a single outer-halo cluster in NGC 5128.  Shortly afterward, \\citet{hol99} obtained similar results for a selection of inner-halo clusters. In this paper, we present new imaging data for another sample of clusters in this important galaxy, more than doubling the total sample.  Throughout this paper we assume $d = 4.0$ Mpc for NGC 5128 and a foreground (Milky Way) reddening of $E_{B-V} = 0.11$, for an apparent distance modulus $(m-M)_V = 28.35$ \\citep{har00}.  At this distance, 1 arcsecond is equivalent to a linear scale of 19.4 pc. Our adopted distance is a mean of the results from three methods including the red-giant-branch tip luminosity \\citep{har99}, the planetary nebula luminosity function \\citep{hui93}, and the $I-$band surface brightness fluctuation technique \\citep{ton01}.  The mean distance is likely to be uncertain to $\\pm0.2$ Mpc based on the close mutual agreement of these methods. ",
        "conclusions": "We have used new imaging data from the HST STIS and WFPC2 cameras to derive structural parameters for globular clusters in the halo of the giant elliptical galaxy NGC 5128. We find that classic, single-mass King models describe their observed light profiles extremely well, allowing us to derive parameters $(r_c, r_h, c, \\epsilon, L)$ for direct comparison with the globular clusters in other galaxies.  The NGC 5128 clusters occupy very much the same regions of parameter space as those in the Milky Way, with the exception that they have a higher range of ellipticities:  they occupy the range $0 < \\epsilon < 0.3$ more or less uniformly, and among various comparison galaxies within the Local Group, we find that they most nearly resemble the old clusters in the LMC and M31 in this respect. We also find half a dozen luminous clusters which may have ``extratidal light'' which is possibly due to active tidal stripping or residual field-star populations from disrupted dwarf satellite galaxies, but may also be the signature of anisotropic velocity distributions.  Lastly, we find that the NGC 5128 clusters delineate a relation between binding energy $E_b$ and luminosity $L$ which is even tighter than in the Milky Way and in exactly the same region of the ``fundamental plane''.  This work provides additional evidence that globular cluster formation processes were remarkably similar in galaxies of very different types.  Considerable further progress can be made in understanding the structures of clusters in this keystone galaxy if we can obtain a more extensive sample of objects over a wide range of galactocentric distances and at the highest resolution possible. In addition, direct spectroscopic measurements of their velocity dispersions are needed to check the key assumptions we have made about their mass-to-light ratios."
    },
    "0206/hep-th0206042_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the most important problems which is facing theoretical physics now is the blending of the Standard Model (SM) with General Relativity (GR). Whatever way we choose (the most popular ones  nowadays are based on some multidimensional constructions involving extended objects), nobody doubts that it will definitely modify physics at {\\it short scales}. On the other hand, the current general paradigm is to keep General Relativity unchanged at {\\it large scales}, but to add new forms of gravitating matter beyond the Standard Model (dark matter, dark energy) for explaining pressing astrophysical and cosmological facts such as galactic rotational curves and the accelerating universe. In the present paper, we consider an alternative paradigm: a {\\it modification} of General Relativity at large scales as a possible explanation of some pressing cosmological issues (notably cosmic acceleration).  The modification of GR that we are going to consider is linked to the issue of ``massive gravity'' (for very light gravitons, with Compton wavelength of cosmological scale). A generic prediction of multidimensional constructions is the existence of massive gravitons. In particular, any Kaluza-Klein (KK) model predicts, besides a massless graviton, the presence of an infinite tower of massive gravitons. However, it seems impossible to use the tower of massive KK gravitons to modify gravity at large scales. Indeed, its spectrum is generically regularly spaced (as illustrated on Fig.~1a), so that, even if the first mode were very light (i.e. of cosmological Compton wavelength), there would exist no regime where the first mode (or first few modes) would be important, and where one could truncate away the rest of the tower of massive states. In other words, as soon as the first mode is important, we open the extra KK dimensions (see, however, below). The situation is, however, different in some brane models. In particular, Refs.~\\cite{Kogan:2000wc}-\\cite{Kogan:2000xc} discovered the possibility (illustrated in Fig.~1b or Fig.~1c) of having a {\\it hierarchical gap}, $m_1 \\ll m_2$, between the first mode (or first group, or even band, of modes) and the tower of higher modes. This situation, called {\\it multigravity} (see \\cite{Kogan:2001ub} for a review and \\cite{Papazoglou:2001cc} for detailed presentation), makes it possible to envisage an effective four-dimensional theory which contains only the massless and ultra-light gravitons and discards the states of mass $m \\geq m_2$. The constructions \\cite{Kogan:2000wc}-\\cite{Kogan:2000xc} predict see-saw-like spectra, $m_1 \\, M_{\\rm Planck}^{1+\\gamma} \\sim m_2^{2+\\gamma}$, with $\\gamma$ interpolating \\cite{Kogan:2001ub} between $0$ \\cite{Kogan:2000wc} and $1$ \\cite{Gregory:2000jc}. Such spectra are naturally compatible with the phenomenologically interesting situation where $m_1^{-1}$ is of cosmological order, while $m_2^{-1}$ is smaller than the millimetre scale. \\begin{figure}[ht] \\begin{center} \\epsfxsize=5in \\epsfbox{spectrum.eps} \\caption{Regular spectrum on (Fig.1 a) versus bigravity (Fig.1 b) or quazi-localized gravity (Fig. 1 c). The last spectrum is continuous but the first band is very narrow in comparison with the gap between bands.} \\end{center} \\end{figure}  So far multigravity was only analyzed in the linearized approximation. The main emphasis of this paper is to provide a fully non-linear formulation of multigravity, i.e. to write down, and analyze, a class of consistent effective four-dimensional Lagrangians, describing, in some limit, the light-mode truncation of the hierarchical spectra of Figs.~1b or 1c. Though we shall illustrate below our approach in the context of particular multidimensional realizations (notably brane models exhibiting multilocalization \\cite{Kogan:2000wc},\\cite{Kogan:2001wp} or  quasi-localization \\cite{Gregory:2000jc}, \\cite{Csaki:2000pp}, \\cite{Dvali:2000rv}), we view our considerations as concerning a very general phenomenon: the concept of {\\it Weakly Coupled Worlds} (WCW). The concept of WCW is very simple: one assumes that there are several Universes (labelled by $i = 1 , \\ldots , N$), each endowed with its own metric $g_{(i)\\mu\\nu}$ and set of matter fields $\\{ \\Phi_i \\}$, which are coupled only through some mixing of their gravitational fields. We require that the theory describing the WCW be near a point of enhanced symmetry, in the sense that there exists a limit (say as some parameter $\\lambda \\rightarrow 0$) where the theory contains $N$ diffeomorphism-like symmetries, corresponding to $N$ massless gravitons. A recent theorem \\cite{BDGH} has proven that the only consistent non-linear theory involving $N$ massless gravitons is the sum of $N$ decoupled GR-type actions \\be \\label{eq1} S_0 = \\sum_{i=1}^{N} S [g_i , \\Phi_i ] \\, , \\ee with (we use the signature $-+++$) \\be \\label{eq2} S [g_i , \\Phi_i ] = \\int d^4 x \\, \\sqrt{-g_i} \\, [M_i^2 R (g_i) - \\Lambda_i + L (g_i , \\Phi_i)] \\, . \\ee Therefore, the only consistent action for a theory of worlds coupled only through gravity is of the form \\be \\label{eq3} S_{\\rm tot} = \\sum_{i=1}^{N} S [g_i , \\Phi_i ] + \\lambda \\, S_{\\rm int} (g_1 , g_2 , \\ldots , g_N) \\, . \\ee When $\\lambda \\rightarrow 0$, the $N$ worlds are non interacting (which implies that, from the point of view of any observer in one world, the other worlds have only a meta-physical existence), and the theory has the enormous symmetry $\\Pi_i {\\rm Diff}_{(i)}$, where each diffeomorphism group ${\\rm Diff}_{(i)}$ acts separately on its own metric $g_{(i)\\mu\\nu}$ and matter fields $\\{ \\Phi_i \\}$. In the interacting case, $\\lambda \\ne 0$, the symmetry of the full action must (again because of the theorem \\cite{BDGH}) be reduced to (at most) one group of diffeomorphisms: the diagonal group of common diffeomorphisms transforming all metrics as \\be \\label{eq4} \\delta \\, g_{\\mu \\nu}^{(i)} = \\epsilon^{\\lambda} \\, \\partial_{\\lambda} \\, g_{\\mu \\nu}^{(i)} + \\partial_{\\mu} \\, \\epsilon^{\\lambda} \\, g_{\\lambda \\nu}^{(i)} + \\partial_{\\nu} \\, \\epsilon^{\\lambda} \\, g_{\\mu \\lambda}^{(i)} \\equiv D_{\\mu}^{(i)} \\, \\epsilon_{\\nu} + D_{\\nu}^{(i)} \\, \\epsilon_{\\mu} \\, . \\ee This symmetry restricts the interaction term $\\lambda \\, S_{\\rm int} (g_1 , \\ldots , g_N)$ to depend only on the invariants one can make with several metrics. This even leaves room for extra kinetic terms built from covariant derivatives such as $g_{(i)}^{\\mu \\nu} \\, D_{\\lambda}^{(j)} \\, g_{\\mu \\nu}^{(k)}$ (such terms do not exist in the case of one metric because $D_{\\lambda}^{(i)} \\, g_{\\mu \\nu}^{(i)} \\equiv 0$). However, in view of the many potential diseases associated to modifications of the standard Einsteinian kinetic terms, and in the spirit of describing the class of interaction terms most relevant at large scales,\\footnote{See Section 2.1 below for further discussion about extra kinetic terms.} i.e. containing the lowest possible number of derivatives ( namely zero, as expected from a generalization of the mass terms that appear in linearized multigravity), we shall only consider ultra-local interaction terms, i.e. \\be \\label{eq5} \\lambda \\, S_{\\rm int} = -\\mu^4 \\int d^4 x \\, {\\cal V} (g_1 (x) , \\ldots , g_N (x)) \\, , \\ee where $\\mu$ is a mass scale (henceforth replacing $\\lambda$ as ``small parameter'') and where ${\\cal V}$ is a scalar density made out of the values of the $N$ metrics at the same ``point''. We assume, for simplicity, that the $N$ weakly coupled worlds ``live'' on the same abstract manifold, i.e., in other terms, that one is given a family of (smooth) canonical one-to-one maps: ${\\rm world}_{(i)} \\rightarrow {\\rm world}_{(j)}$.  The aim of this paper is threefold: (i) to motivate the possibility of the effective action (\\ref{eq3}), (\\ref{eq5}) by considering several different specific models (brane models, Kaluza-Klein models and non-commutative geometry ideas); (ii) to delineate and parametrize the various ``universality classes'' of non-linear multigravity; and (iii) to sketch the main qualitative consequences of such non-linear multigravity theories and to contrast them with the usual paradigm of ``massless plus massive gravitons'' which is based on a linearized approximation.  It should be noted that theories defined by (\\ref{eq3}), (\\ref{eq5}) (in the ``bigravity'' case: $N=2$) were first introduced in the seventies \\cite{ISS} as a model for describing a sector of hadronic physics where a massive spin-2 field (the ``$f$ meson'', with ``Planck mass'' $M_f \\sim 1 \\, {\\rm GeV}$ in Eq.~(\\ref{eq2})) plays a dominant role. It was then called ``strong gravity'' or the ``$f$-$g$ theory''. Our work not only proposes to revive, within a new (purely ``gravitational'') physical context, this early proposal, but initiates the task of systematically studying the general phenomenological consequences of the action (\\ref{eq3}), (\\ref{eq5}). The present paper will only briefly sketch the new physical paradigm following from such actions. In subsequent papers, we shall discuss in detail the cosmological consequences of such theories \\cite{DKP1}, as well as its strong-field phenomenology \\cite{DKP2}. ",
        "conclusions": "In this paper we suggested a new  paradigm concerning ``massive gravity'' and ``large scale modification of gravity''.  Considering the fully nonlinear  bigravity action suggests to change viewpoint: instead of the theory with massless and massive graviton(s) we had in linearized approximation, we are dealing with  several interacting metrics. We introduced the  concept of universality class which we formulated using bigravity (two interacting metrics) as an example. Different approaches (brane, KK, non-commutative geometry) naturally lead to different universality classes for the fully nonlinear bigravity action. Another important new  suggestion  is that almost all solutions must  now be  of the  non-asymptotically flat (cosmological) type.   This new formulation can change the standard problematic of the  $m^2 \\rightarrow 0$ discontinuity. We showed the existence of classes of solutions that are compatible with ``our universe''. However, we do not claim to have proven that general solutions of bigravity are phenomenologically acceptable. The two main problems of massive gravity (ghost, potential blow up of some field variables when $m^2 \\rightarrow 0$) must still be examined in detail. The  important  problem is to find  the  matching to the local sources of the field so that the full metric is free of singularities. We do not worry about matching at infinity because we abandon the requirement of asymptotic flatness. It is possible that in some models of bigravity such local matching does not exist  because of the explicit or implicit presence of ghost modes in the theory.  Such models  would be  physically unacceptable. We note in this respect that  the 6-dimensional model discussed in  \\cite{Kogan:2001yr} which does not  contain negative tension branes, contains instead either  branes with equations of state violating the weak energy condition $T_{\\mu\\nu}^{\\rm brane} \\, \\ell^{\\mu} \\, \\ell^{\\nu} \\geq 0$ ( with light-like $\\ell^{\\mu}$) or has a conifold singularity in the bulk. The physical consistency of this model must be further investigated.  We have also quoted mathematical theorems linking the existence of  a hierarchical  spectrum (necessary  for the derivation of an effective bigravity Lagrangian) to the necessary negativity of the Ricci curvature of the compactified manifold. This sign condition might hide the presence of ghost-like fields in the theory. These questions are  pressing and  deserve detailed investigation.  Assuming a positive resolution of these issues or simply taking the phenomenological viewpoint that nonlinear bigravity Lagrangians open an interesting new arena for non standard gravitational effects, we shall explore in future publications \\cite{DKP1}, \\cite{DKP2} the nonlinear physics of bigravity actions, with a particular view on its cosmological aspects, as it may provide a natural candidate for some new type of ``dark energy''.   \\vskip1cm  \\textbf{Acknowledgments:} We would like to thank P. B\\'{e}rard, M. Berger, A. Connes, J. Fr\\\"{o}hlich, M. Gromov, M. Kontsevitch, A. Papazoglou, G. Ross and A. Vainshtein  for  informative  discussions. I.K. is supported in part by PPARC rolling grant PPA/G/O/1998/00567 and  EC TMR grants HPRN-CT-2000-00152 and  HRRN-CT-2000-00148."
    },
    "0206/astro-ph0206019_arXiv.txt": {
        "abstract": "We studied the relation between the ratio of rotational velocity to velocity dispersion and the metallicity (\\vsigmr ) of globular cluster systems (GCS) of disk galaxies by comparing the relation predicted from simple chemo-dynamical models for the formation and evolution of disk galaxies with the observed kinematical and chemical properties of their GCSs.  We conclude that proto disk galaxies underwent a slow initial collapse that was followed by a rapid contraction and derive that the ratio of the initial collapse time scale to the active star formation time scale is $\\sim 6$ for our Galaxy and $\\sim 15$ for M31.  The fundamental formation process of disk galaxies was simulated based on simple chemo-dynamical models assuming the conservation of their angular momentum. We suggest that there is a typical universal pattern in the \\vsigmr \\ of the GCS of disk galaxies. This picture is supported by the observed properties of GCSs in the Galaxy and in M31. This relation would deviate from the universal pattern, however, if large-scale merging events took major role in chemo-dynamical evolution of galaxies and will reflect the epoch of such merging events. We discuss the properties of the GCS of M81 and suggest the presence of past major merging event. ",
        "introduction": "The importance of investigations of globular clusters (GCs) has been reassessd in recent years (Ashman \\& Zepf 1998, Kissler-Patig 2000), particularly since the globular cluster systems (GCS) of disk galaxies can be used to study the dynamical formation history of their host galaxies.  The rapid dynamic formation scenario of our Galaxy was first proposed by Eggen, Lynden-Bell \\& Sandage (1962; hereafter ELS) based on apparent correlations among the kinematic properties and metallicity indicators of the halo and disk stars in the solar neighborhood.  Previous studies interpreted the properties of the GCS in relation to the formation of our Galaxy. Maline, Hartmann \\& Mathews (1991) treated the GCs of the Galaxy as a tracer of its chemical evolution and concluded that the Galaxy was formed by an inhomogeneous collapse.  Kinman (1959) demonstrated that the GCS of our Galaxy as a whole has a small net rotation and Frenk \\& White (1980) suggested that the difference in rotational velocity between the F-type (metal-poor) and G-type (metal-rich) GC groups is only marginally significant in our Galaxy. Huchra et al. (1982) also discovered a rotating system of M31 GCs. Zinn (1993) pointed out that inner \"old halo\" clusters have rapid rotation, small velocity dispersion, and flattened orbits and suggested that these properties are evidence of the dissipational collapse formation of our Galaxy.  Minniti (1996) studied the kinematics of GCs and stars in the Galactic bulge and concluded that the metal-rich GCs and stars in the Galactic bulge have similar kinematic properties to the inner \"old halo\" population (cf. Zinn 1993) and that the Galactic bulge was formed by dissipational collapse.  Recently, Chiba \\& Beers (2000) discussed a new scenario for the formation of our Galaxy from studying the proper motions of metal-poor halo stars in the solar neighborhood obtained by HIPPARCOS.  In studying the global dynamic evolution of galaxies, it is advantageous to use the kinematic properties of the GCSs of those galaxies, because most GCs should have been formed before, during, or immediately after the collapse of their host galaxies (Fall \\& Rees 1985). Those GCs, therefore, can give us information on the dynamic evolution that occurred during those epochs.  We discuss the dynamic formation scenario of disk galaxies using the kinematic property of the GCS, especially the relation between metallicity and the ratio of the rotational velocity to the velocity dispersion.  We construct simple models to simulate the evolution of the chemo-dynamic properties of the GCSs of disk galaxies in Section 2 and compare the observed properties of the GCSs of our Galaxy and M31 with those derived from these simple models in Section 3. In Section 4, we discuss the importance of observing GCSs in S0 and Sa type galaxies. Finally, in Section 5, we try to recompile the GCS data of M81 and propose that the chemo-dynamic evolution of the M81 system may differ from those of the Galaxy and M31. We then conclude that we must investigate the detailed kinematics of the M81 GCS, because the quality and quantity of the data on M81 GCs is poor. ",
        "conclusions": "We made a series of chemo-dynamic model calculations to predict the relation for GCSs under the assumption that GCSs are formed before or during the collapse of protogalactic gas clouds.  Three collapse profiles, e.g., exponential collapse, delayed collapse, and early collapse, are assumed in our simple models. The power index $n$ of the Schmidt Law for the star formation process, the dynamic collapse time scale $t_{col}$, and the star formation time scale $t_{sf}$ are the free parameters chosen in this study.  We analyzed the \\vsigmr s of our Galaxy (Figure 4) and M31 (Figure 5). The observed relations for these two disk galaxies of similar mass and morphological type are similar. They exhibit a rather sudden increase in \\vsig\\ at [Fe/H]$\\sim -1.0$ for the Galaxy and at [Fe/H]$\\sim -0.6$ for M31. This is consistent with chemo-dynamic model predictions of the formation of the GCSs of these galaxies. However we have to notice that it is difficult to choose the parameter set ($t_{col,1}$,$t_{col,2}$) for fitting the model with data points for early collapse model. It also suggests that the GCSs of these galaxies were formed in a dynamic collapse, where the collapse time scale is as large as, or larger than, the star formation time scale, if the star formation process is proportional to the protogalactic gas density, namely $n=1$. The collapse time scale should be at least an order of magnitude larger than the star formation time scale, if the star formation process took place more efficiently in an enhanced density region, as is the case for $n=2$.  We noticed a difference in the [Fe/H] value at which the abrupt increase of \\vsig\\ takes place between the Galaxy and M31. By comparing these observed curves with the results of simple models, we derived ratios of the collapse time scale to the star formation time scale for the Galaxy and M31 of $t_{col}/t_{sf} \\sim 3-11$ and $t_{col}/t_{sf} \\geq 7$, respectively, when the power index of the Schmidt Law for star formation is $n=1$. For $n=2$, those ratios are $t_{col}/t_{sf} \\geq 100$. This might be a evidence that the $n=2$ model is ruled out, however we have to construct more realistic chemo-dynamical evolution model for discussing whether it is ruled out or not. These differences in the [Fe/H] value do not reflect a significant differences in the $t_{col}/t_{sf}$. We suggest that the difference of the \\vsigmr \\ between galaxies should be a probe for studying the differences of formation scenarios of those galaxies, although we cannot restrict the formation scenario of a galaxy from it.   For the GCSs of both our Galaxy and M31, we see that the value of \\vsig\\ is $\\sim 0$ at $-1.5<$[Fe/H]$<-0.8$, which is not consistent with the models. Since we assumed that the protogalactic gas was monolithic in our models, \\vsig\\ may have been affected by merger events during the epoch that correspond to that metal range.  GCs are useful probes for investigating the chemo-dynamic properties of host galaxies. New observations of the GCSs of nearby disk galaxies with 8-10m class telescopes will provide fruitful suggestions on the history of galaxy formation.  \\appendix"
    },
    "0206/astro-ph0206475_arXiv.txt": {
        "abstract": "{We analyse the optical and IR spectra, as well as the spectral energy distribution (UV to mm) of the candidate Herbig Ae star HD~100453. This star is particular, as it shows an energy distribution similar to that of other isolated Herbig Ae/Be stars (HAEBEs), but unlike most of them, it does not have a silicate emission feature at 10~\\mic, as is shown in \\citet{meeus2001}. We confirm the HAEBE nature of HD~100453 through an analysis of its optical spectrum and derived location in the H-R diagram. The IR spectrum of HD~100453 is modelled by an optically thin radiative transfer code, from which we derive constraints on the composition, grain-size and temperature distribution of the circumstellar dust. We show that it is both possible to explain the lack of the silicate feature as (1) a grain-size effect - lack of {\\bf small} silicate grains, and (2) a temperature effect - lack of small, {\\bf hot} silicates, as proposed by \\citet{dullemond2001}, and discuss both possibilities. \\cdrm{Finally, we show that the latter possibility is the more preferable.} ",
        "introduction": "Herbig Ae/Be stars are intermediate-mass pre-main sequence stars which show an IR excess due to circumstellar (CS) dust. This dust is believed to be located in a disc \\citep[see e.g.][]{adams1987ApJ...312..788A,beckwith1990AJ.....99..924B,meeus1998A&A...329..131M}; \\citet{chiang1997ApJ...490..368C} were the first to develop a CS disc model in which a warm layer of optically thin material surrounds the (cold) midplane. This passive disc model provides an elegant solution to the puzzling observation that the disc must be optically thick in the near-IR and mid-IR (in order to account for the sub-millimeter flux) while the Si-O stretch of amorphous silicates at 9.7 $\\mu$m is observed to be in emission (implying it is caused by optically thin dust).  Since small amorphous silicate grains are the most abundant dust species in interstellar space, their spectral signature is expected to be present during virtually all stages of the star formation process: in absorption during the protostellar collapse and active accretion disc phases, and in emission during the passive disc and debris disc phases. Indeed, observations of star forming regions confirm this picture. However, during our analysis of the infrared spectra of passive discs surrounding isolated Herbig Ae/Be stars taken with the Infrared Space Observatory \\citep[ISO;][]{kessler1996A&A...315L..27K}, we found that some stars do not show any evidence for the presence of the 9.7 $\\mu$m band \\citep[][ hereafter Paper~I]{meeus2001}. In Paper~I we speculate that the lack of silicate emission is due to the absence of small silicate grains in the disc, due to grain growth and the removal of small grains by radiation pressure.  \\citet[][ hereafter DDN]{dullemond2001} propose a different explanation, by avoiding the presence of small silicate grains in a certain temperature range. In the DDN model, the disc inner rim has a large scale-height due to the fact that it receives direct stellar radiation (an effect not taken into account in the Chiang \\& Goldreich models). The inner rim causes a shadowed region behind the rim where the temperature is low; at some distance from the star the surface of the flaring disc emerges from the shadow and receives direct starlight. The rim causes two effects: (1) it creates a prominent near-IR flux contribution which is in good agreement with observations \\citep[see][ DDN]{natta2001A&A...371..186N}, and (2) for certain rim heights the shadowed region can suppress the strength of the 9.7 $\\mu$m silicate emission (see DDN).  Stimulated by these considerations, we decided to derive empirical constraints on the amount of mass that can be present in the form of warm optically thin silicate grains in HAEBE stars that lack the silicate feature. We investigate the effect the size and the temperature distribution of dust grains can have on the appearance of spectral features. Therefore, we look for observational constraints on the average size of the silicate grains and the maximal mass of small, warm silicate grains. This we do by modelling the isolated Herbig Ae star HD~100453, which is one of the four isolated HAEBEs in our sample where the silicate feature is found to be absent \\citep{meeus2001}. We selected this star because it is the brightest of all four, resulting in observations with higher signal to noise ratios. The model we used to derive the constraints is the optically thin radiative transfer code MODUST (Bouwman \\& de Koter, in preparation). ",
        "conclusions": "We searched for an explanation for the absence of the 10~\\mic \\ silicate feature in the spectrum of HD~100453 in terms of either a size or a temperature effect. We first showed, from a study of the optical spectra, that the star belongs to the group of the isolated HAEBEs. HD~100453 is member of group Ib, which contains stars with an increasing IR SED that lack the silicate emission feature (see paper I). Our conclusions can be summarised as follows:  \\begin{enumerate}  \\item{we confirmed the HAEBE nature of HD~100453 with the observation of the \\halfa \\ line in emission, as well as the presence of other typical CS HAEBE lines; furthermore, the star has an age of approximately 10 Myr and is located close to the ZAMS}  \\item{we modelled the IR spectrum of HD~100453 with an optically thin radiative transfer code, with the following properties:  \\begin{itemize}  \\item the dust emission stems from two different temperature regimes, with an average temperature of $\\sim$ 360 K and $\\sim$ 60 K; the bulk of the mass consists of amorphous silicates, but carbonaceous material and iron(oxide) are present as well  \\item the absence of the silicate emission feature can be caused by the absence of {\\em small} particles: a size distribution starting at 4~\\mic, with an average of 6.7~\\mic \\ does not reveal the silicate feature. The minimum average radius derived for the silicate particles is about 355 times larger than that for AB Aur, suggesting that, in the disc of HD~100453, grain growth has taken place on a much larger scale than in the disc of AB Aur.  \\item we exclude dust removal processes such as radiation pressure and P-R drag to be the main cause for the difference in grain size, given the larger luminosity of AB Aur and, more important the presence of CS gas which makes this processes inefficient  \\item the presence of small (hot) particles of species other than silicates (such as PAHs and metallic iron) is established. Furthermore, from a comparison of the slope at longer (IR-mm) wavelengths with other HAEBEs showing the silicate feature, we deduce a lack of large grains in the CS disc of HD~100453. \\cdrev{These may be seen as arguments against the coagulation explanation for the absence of the silicate feature.}  \\end{itemize} }  \\item{alternatively, the absence of the silicate emission feature can be caused by the absence of a large amount of small, hot silicates; colder (T $<$ 200 K), small silicates still can be present in a much larger amount. We obtained a maximal mass of small, hot silicates between 10$^{-12}$ and 10$^{-11}$ \\Msun, which is 0.2 to 0.7 \\% of the hot dust mass, and roughly a factor 10$^{-6}$ of the total dust mass. When comparing HD~100453 with AB Aur, the maximal mass residing in the small, hot grains is a factor hundred to five hundred smaller. The derived maximal mass is in agreement with the theoretical predictions by DDN. They propose shielding by a puffed-up inner region as an attractive explanation for the absence of hot, small silicates; and hence for the absence of the 10 micron silicate emission feature.}  \\end{enumerate}  \\noindent High spatial resolution observations are needed to give constraints on the location of the grains emitting in the infrared; also more long-wavelength observations are necessary to study the coldest dust. Only then will we be able to clearly distinguish between both effects."
    },
    "0206/astro-ph0206196_arXiv.txt": {
        "abstract": "The intense turbulence present in the solar convection zone is a major challenge to both theory and simulation as one tries to understand the origins of the striking differential rotation profile with radius and latitude that has been revealed by helioseismology.  The differential rotation must be an essential element in the operation of the solar magnetic dynamo and its cycles of activity, yet there are many aspects of the interplay between convection, rotation and magnetic fields that are still unclear.  We have here carried out a series of 3--D numerical simulations of turbulent convection within deep spherical shells using our anelastic spherical harmonic (ASH) code on massively parallel supercomputers.  These studies of the global dynamics of the solar convection zone concentrate on how the differential rotation and meridional circulation are established.  We have addressed two issues raised by previous simulations with ASH.  Firstly, can solutions be obtained which possess the apparent solar property that the angular velocity $\\Omega$ continues to decrease significantly with latitude as the pole is approached?  Prior simulations had most of their rotational slowing with latitude confined to the interval from the equator to about 45$^{\\circ}$.  Secondly, can a strong latitudinal angular velocity contrast $\\Delta \\Omega$ be sustained as the convection becomes increasingly more complex and turbulent?  There was a tendency for $\\Delta \\Omega$ to diminish in some of the turbulent solutions that also required the emerging energy flux to be invariant with latitude.  In responding to these questions, five cases of increasingly turbulent convection coupled with rotation have been studied along two paths in parameter space.  We have achieved in one case the slow pole behavior comparable to that deduced from helioseismology, and have retained in our more turbulent simulations a consistently strong $\\Delta \\Omega$. We have analyzed the transport of angular momentum in establishing such differential rotation, and clarified the roles played by Reynolds stresses and the meridional circulation in this process.  We have found that the Reynolds stresses are crucial in transporting angular momentum toward the equator.  The effects of baroclinicity (thermal wind) have been found to have a modest role in the resulting mean zonal flows. The simulations have produced differential rotation profiles within the bulk of the convection zone that make reasonable contact with ones inferred from helioseismic inversions, namely possessing a fast equator, an angular velocity difference of about 30\\% from equator to pole, and some constancy along radial lines at mid-latitudes.  Future studies must address the implications of the tachocline at the base of the convection zone, and the near-surface shear layer, upon that differential rotation. ",
        "introduction": "The solar turbulent convection zone has striking dynamical properties that continue to challenge basic theory.  The most fundamental issues involve the solar rotation profile with latitude and depth, and the manner in which the 22-year cycles of solar magnetic activity are achieved.  These two issues are closely interrelated, for the global dynamo action is likely to be very sensitive to the angular velocity $\\Omega$ profiles realized by convection redistributing angular momentum within the deep zone.  Both dynamical topics touch on the seeming inconsistency that turbulence can be both highly intermittent and chaotic on smaller spatial and temporal scales, yet exhibit large-scale ordered behavior (cf.  Brummell, Cattaneo \\& Toomre 1995). The differential rotation profile established by the turbulent convection, though strong in contrast, is remarkably smooth; the global-scale magnetic activity is orderly, involving sunspot eruptions with very well defined rules for field parity and emergence latitudes as the cycle evolves.  The wide range of dynamical scales of turbulence present in the solar convection zone yield severe challenges to both theory and simulation:  the discernible structures range from granules ($\\sim 10^3$ km or 1 Mm in horizontal size), to supergranules ($\\sim$30 Mm), to possible patterns of giant cells comparable to the overall depth of that zone ($\\sim$200 Mm, or nearly 30\\% by radius).  Given that the dissipation scales are on the order of 0.1 km or smaller, the solar turbulence encompasses at least six orders of magnitude for each of the three physical dimensions.  The largest current 3--D turbulence simulations can resolve about three orders of magnitude in each dimension.  Yet despite the vast difference in the range of scales dynamically active in the sun and those accessible to simulations, the latter have begun to reveal basic self-ordering dynamical processes yielding coherent structures that appear to play a crucial role in the global differential rotation and magnetic dynamo activity realized in the sun.  It has long been known by tracking surface features that the surface of the sun rotates differentially (e.g. Ward 1966, Sch\\\"ussler 1987):  there is a smooth poleward decline in the angular velocity $\\Omega$, the rotation period being about 25 days in equatorial regions and about 33 days near the poles.  Helioseismology, which involves the study of the acoustic $p$-mode oscillations of the solar interior (e.g. Gough \\& Toomre 1991), has provided a remarkable new window for studying dynamical processes deep within the sun.  This has been enabled by nearly continuous helioseismic observations provided from the SOHO spacecraft with the high-resolution Michelson Doppler Imager (SOI--MDI) (Scherrer et al. 1995) and from the ground-based Global Oscillation Network Group (GONG) set of six related instruments (Harvey et al. 1996).  The helioseismic findings about differential rotation deeper within the sun have turned out to be revolutionary, for they are unlike any anticipated by convection theory prior to such probing of the interior of a star.  Helioseismology has revealed that the rotation profiles obtained by inversion of frequency splittings of the $p$ modes (e.g. Libbrecht 1989, Thompson et al. 1996, Schou et al. 1998, Howe et al. 2000b) have the striking behavior shown in Figure \\ref{fig1}.  The variation of angular velocity $\\Omega$ observed near the surface, where the rotation is considerably faster at the equator than near the poles, extends through much of the convection zone with relatively little radial dependence.  Thus at mid-latitudes $\\Omega$ is nearly constant on radial lines, in sharp contrast to early numerical simulations of rotating convection in spherical shells (e.g. Gilman \\& Miller 1981, Glatzmaier 1987) that suggested that $\\Omega$ should be nearly constant on cylinders aligned with the rotation axis and decreasing inward on the equatorial plane. Another striking feature is the region of strong shear at the base of the convection zone, now known as the tachocline, where $\\Omega$ adjusts to apparent solid body rotation in the deeper radiative interior.  Whereas the convection zone exhibits prominent differential rotation, the deeper radiative interior does not; these two regions are joined by the complex shear of the tachocline.  There is further a thin shear boundary layer near the surface in which $\\Omega$ increases with depth at intermediate and high latitudes.  \\begin{figure*}[!htp] \\centerline{\\includegraphics[width=0.95\\linewidth]{./f1.eps}}  \\caption[]{\\label{fig1} ~($a$) Angular velocity profile $\\Omega/2\\pi$ with radius and latitude as deduced from helioseismology using SOI--MDI data, with red tones indicating fast rotation and blue-green the slowest rotation [adapted from Schou et al. 1998].  ~($b$) Time-averaged rotation rates from five years of GONG helioseismic data, plotted against radius at different latitudes. The surface shear layer and the tachocline at the base of the convective zone are indicated, as well as the zone covered by our computational domain (grey area) [adapted from Howe et al. 2000b].} \\end{figure*}  The tachocline has been one of the most surprising discoveries of helioseismology, especially since its strong rotational shear affords a promising site for the solar global dynamo.  Such a tachocline was not anticipated, and current theoretical approaches to explain its presence are still only innovative sketches (Spiegel \\& Zahn 1992; Gough \\& McIntyre 1998; Charbonneau, Dikpati \\& Gilman 1999).  Helioseismology has also recently detected prominent variations in the rotation rate near the base of the convective envelope, with a period of 1.3 years evident at low latitudes (Howe et al. 2000a; Toomre et al. 2000). These are the first indications of dynamical changes close to the presumed site of the global dynamo as the cycle advances.  Such a succession of developments from helioseismology provide both a challenge and a stimulus to theoretical work on solar convection zone dynamics.  Seeking to understand solar differential rotation and magnetism requires 3--D simulations of convection in the correct full spherical geometry.  However, the global nature of such solutions represent a major computational problem, given that the largest scale is pinned and only three orders of magnitude smaller in scale can be represented. Much of the small-scale dynamics in the sun dealing with supergranulation and granulation are by necessity then largely omitted.  The alternative is to reduce the fixed maximum scale by studying smaller localized domains within the full shell and utilizing the three orders of magnitude to encompass the dynamical range of turbulent scales.  There are clear tradeoffs: ~the global models operate in the correct geometry yet struggle to encompass enough of a dynamical range to admit fully turbulent solutions, whereas the local models are able to study intensely turbulent convection but only within a particular limited portion of the full domain.  Both approaches are needed, and the efforts are complementary, as reviewed in detail by Gilman (2000) and Miesch (2000).  Highly turbulent but localized 3--D portions of a convecting spherical shells are being studied to assess transport properties and topologies of dynamical structures (e.g. Brandenburg et al. 1996; Brummell et al. 1996, 1998; Porter \\& Woodward 2000; Robinson \\& Chan 2001), of penetration into stable domains below (Brummell et al. 2001, Porter \\& Woodward 2001), of effects of realistic near-surface physics on granulation and supergranulation (e.g. Stein \\& Nordlund 1998), and of dynamo processes and magnetic transport by the convection (e.g. Cattaneo 1999, Tobias et al.  2001). Without recourse to direct simulations, the angular momentum and energy transport properties of turbulent convection have also been considered using mean-field approaches to derive second-order correlations (the Reynolds stresses and anisotropic heat transport) under the assumption of separability of scales.   Although such procedures involve major uncertainties, the resulting angular momentum transport, which is described by mechanisms such as the so-called $\\Lambda$ effect, have served to reproduce the solar meridional circulation (e.g. Durney 1999, 2000) and differential rotation (e.g.  Kichatinov \\& R\\\"{u}diger 1995). Various other states can be achieved by adjusting parameters.  Initial studies of convection in full spherical shells to assess effects of rotation with correct account of geometry (e.g. Gilman \\& Miller 1981; Glatzmaier \\& Gilman 1982; Glatzmaier 1985, 1987; Sun \\& Schubert 1995) have set the stage for our efforts to study more turbulent flows using new numerical codes designed for the massively parallel computer architectures that are enabling such major simulations.  We here report on our continuing studies with the anelastic spherical harmonic (ASH) code (Clune et al.  1999) to examine the $\\Omega$ profiles established within the bulk of the solar convection zone by turbulent convection, building on the progenitor work by Miesch et al. (2000), Elliott, Miesch \\& Toomre (2000), and Brun \\& Toomre (2001). We also recognize the recent modelling of convection in  spherical shells by Takehiro \\& Hayashi (1999) and Grote \\& Busse (2001).  The simulations reported in Miesch et al. (2000) and Elliott et al. (2000) have revealed the richness and complexity of compressible convection achieved in rotating spherical shells.  Most of the resulting angular velocity profiles in the seven simulations considered have begun to make substantial contact with the helioseismic deductions within the bulk of the solar convection zone. These possess fast equatorial rotation (prograde), substantial $\\Omega$ contrasts with latitude, and reduced tendencies for rotation to be constant on cylinders. The simulations with ASH have not yet sought to deal with questions of the near-surface rotational shear layer nor with the formation of a tachocline near the base of the convection zone. These studies have revealed that to achieve fast equators it is essential that parameter ranges be considered in which the convection senses strongly the effects of rotation, which translates into having a convective Rossby number less than unity for large Taylor numbers.  Such rotationally-constrained convection exhibits downflowing plumes that are tilted away from the local radial direction, resulting in velocity correlations and thus Reynolds stresses that are found to have a significant role in the redistribution of angular momentum.  This seems to provide paths to realize solar-like $\\Omega$ profiles.  Further, it is desirable to impose thermal boundary conditions at the top of the domain that enforce the constancy of emerging flux with latitude in order to be consistent with what appears to be observed.  We wish to focus on two outstanding issues raised by the prior simulations with ASH that need particular attention concerning the differential rotation established within the bulk of the solar convection zone.  As {\\sl Issue 1}, the helioseismic inferences in Figure 1 emphasize that $\\Omega$ in the sun appears to decrease significantly with latitude even at mid and high latitudes, a property which has been difficult to attain in the prior seven simulations.  The substantial latitudinal decrease in angular velocity, say $\\Delta \\Omega$, in the models is primarily achieved in going from the equator to about 45$^\\circ$, with little further decrease in $\\Omega$ achieved at higher latitudes in most of the cases.  Whereas the overall latitudinal contrasts from equator to pole in the models and in the sun are roughly of the same order, the angular velocity in the sun continues to slow down much more as the pole is approached.  Two models, designated as {\\sl LAM} (in Miesch et al. 2000) and {\\sl L3} (in Elliott et al. 2000), do exhibit $\\Omega$ which decrease at high latitudes, but {\\sl LAM} involves an emerging heat flux that varies too much with latitude due to choice of boundary conditions, and {\\sl L3} has an overall $\\Delta \\Omega$ that is only two-thirds of the helioseismic value.  Thus in confronting {\\sl Issue 1}, we will search in parameter space for solutions that can achieve $\\Omega$ profiles in which the decrease with latitude does not taper off at mid latitudes and for which the contrast $\\Delta \\Omega$ is at least comparable to the helioseismic findings.  As {\\sl Issue 2}, with the convection becoming more turbulent, achieved by decreasing either the thermal or viscous diffusivities, there is a tendency for the latitudinal contrast $\\Delta \\Omega$ in the solutions to diminish or even decrease very prominently, thus being at variance with $\\Delta \\Omega$ deduced from helioseismology.  This behavior appears to arise from increasing complexity leading to a weakening of nonlinear velocity correlations that have a crucial role in angular momentum redistribution.  These Reynolds stress terms are strong in the laminar solutions that involve tilted columnar convection cells (`banana cells') aligned with the rotation axis; they weaken as the flows become more intricate, but would be expected to become again significant once coherent structures develop at higher levels of turbulence.  For example, the model {\\sl TUR} (in Miesch et al. 2000) exhibits the emergence of downflow networks involving fairly persistent plumes which possess some of the expected attributes of the coherent structure seen in localized domains of highly turbulent convection (e.g. Brummell et al. 1998).  As a result, {\\sl TUR} has a fairly interesting angular momentum transport attributed to the nonlinear correlations that sustain a level of differential rotation slightly weaker than {\\sl LAM}, but it too has a considerable variation of heat flux with latitude.  The model {\\sl T2} (in Elliott et al. 2000) sought to correct the latter by using modified thermal boundary conditions but appears to not have attained high enough turbulence levels to realize strong coherent structures.  Absent those features, {\\sl T2} yielded $\\Omega$ profiles with a small $\\Delta \\Omega$, and even a slightly slower equatorial rotation rate than that in the mid latitudes.  Thus in confronting {\\sl Issue 2}, we seek turbulent solutions that possess $\\Omega$ profiles with fast equators and strong latitudinal contrasts $\\Delta \\Omega$, and emerging heat fluxes that vary little with latitude.  To achieve this we have considered two paths in parameter space that yield more turbulent solutions by either varying the Prandtl number or keeping it fixed, while maintaining the same rotational constraint as measured by a convective Rossby number.   We describe briefly in \\S2 the ASH code and the set of parameters used for the simulations studied here.  In \\S3 we discuss the properties of rotating turbulent convection and the resulting differential rotation and the meridional circulation for the five cases $A$, $AB$, $B$, $C$ and $D$.  In \\S4 we analyze the transport of angular momentum by several processes and the influence of baroclinic effects in establishing the mean flows.  In \\S5 we reflect on the significance of our findings, and especially in terms of dealing with the two issues raised by the prior simulations with ASH. ",
        "conclusions": "Our five simulations studying the coupling of turbulent convection and rotation within full spherical shells have revealed that strong differential rotation contrasts can be achieved for a range of parameter values.  With these new models, we have focused on two fundamental issues raised in comparing the solar differential rotation deduced from helioseismology with the profiles achieved in the prior 3--D simulations of turbulent convection with the ASH code (Miesch et al. 2000, Elliott et al. 2000).  As {\\sl Issue 1}, the sun appears to possess remarkably slow poles, with $\\Omega$ decreasing steadily with latitude even at mid and high latitudes (Fig. \\ref{fig1}).  In contrast, the previous models showed little variation in $\\Omega$ at the higher latitudes, having achieved most of their latitudinal angular velocity contrast $\\Delta \\Omega$ in going from the equator to about 45$^\\circ$.  As {\\sl Issue 2}, there was a tendency for $\\Delta \\Omega$ to diminish or even decrease sharply within the prior simulations as the convection became more turbulent, yielding values of $\\Delta \\Omega$ that were becoming small compared to the helioseismic deductions.  In seeking to resolve these two issues, we have explored two paths in parameter space that yield complex and turbulent states of convection.  {\\sl Path 1} involves decreasing the Prandtl number in the sequence of cases $A$, $B$ and $C$, while keeping the P\\'eclet number nearly constant.  {\\sl Path 2} maintains a constant Prandtl number as both the Reynolds and P\\'eclet number are increased in the sequence of cases $AB$, $B$ and $D$.  On both paths the convective Rossby number has been chosen to be less than unity, thereby maintaining a strong rotational influence on the convection even as the flows become more intricate.  In dealing with {\\sl Issue 1}, our case $AB$ provides the first indications that it is possible to attain solutions in which the polar regions rotate significantly slower than the mid latitudes (Fig. \\ref{fig4}). There is a monotonic decrease from the fast (prograde) equatorial rate in $\\Omega$ to the slow (retrograde) rate of the polar regions. Further, that case $AB$ has $\\Omega$ nearly constant on radial lines at the higher latitudes, again in the spirit of the helioseismic inferences.  We do not fully understand why in case $AB$ such a strikingly different $\\Omega$ profile results, compared to that in our other solutions (and of the progenitor simulations) in which the contrast $\\Delta \\Omega$ is mainly achieved in the lower latitudes. Our principal clues come from Figure \\ref{fig11} where we find that only in case $AB$ is the Reynolds stress component of the net radial angular momentum flux ${\\cal F}_{r,R}$ (through shells at various radii) uniformly directed outward.  From having examined in detail angular momentum flux streamfunctions (not shown) with radius and latitude consistent with equations (7--9), we observed that the Reynolds stress contributions to such transport possessed multi-celled structures with radius at high latitudes in all the cases except $AB$.  The single-cell behavior there for case $AB$ appears to enable more effective extraction of angular momentum by Reynolds stresses from the high to the low latitudes, thereby yielding a distinctive rotational slowing of the high latitudes. Further, case $AB$ possesses strong meridional circulations at low latitudes, but only feeble ones at latitudes above 30$^{\\circ}$, unlike other solutions such as case $C$ (Figs. \\ref{fig8}, \\ref{fig9}).  This yields a weak meridional component ${\\cal F}_{\\theta,M}$ (seeking to spin up the poles) to the latitudinal angular momentum flux, thereby allowing the equatorward transport by the Reynolds stress component ${\\cal F}_{\\theta,R}$ to succeed in extracting angular momentum from the higher latitudes. Such polar slowing also leads to case $AB$ possessing the greatest $\\Delta \\Omega$ attained in our five simulations (Table $2$).  We also considered the possibility that the slow pole behavior in case $AB$ may have baroclinic origins.  This can result from suitable correlations in velocity and thermal structures yielding a latitudinal heat flux which may produce substantial entropy variations at the higher latitudes, thereby leading to greater baroclinic contributions in equation (11) that defines the variation of mean zonal velocity. Examination of Figure \\ref{fig12} at high latitudes does not reveal a prominent baroclinic contribution, and this is consistent with the bland variation of entropy for case $AB$ (Fig. \\ref{fig14}) at latitudes above about 40$^\\circ$.  We conclude that the origin of the slow rotation rate in polar regions appears to be primarily dynamical, being associated with the Reynolds stress transports, and not with baroclinicity that arises from latitudinal heat transport serving to establish a sufficiently warm pole.  Although case $AB$ provides a solution that resolves {\\sl Issue 1}, it is unique in achieving this among our five simulations.  It may be that in parameter space there only exists a small basin of attraction for such behavior, though we think it more likely that several solution states may coexist for the same control parameters, one of which exhibits the gradual rotational slowing at high latitudes, and others having most $\\Omega$ variations confined to low and mid latitudes.  We plan to examine whether the slow pole characteristics of case $AB$ can be maintained at nearby sites in parameter space if started from initial conditions corresponding to $AB$, and plan to report on this in the future.   {\\sl Issue 2} concerns sustaining a strong differential rotation with latitude as the convection becomes more complex.  The two paths that we have explored in parameter space to achieve more complex and turbulent states yield relative angular velocity contrasts $\\Delta \\Omega/ \\Omega_o$ in latitude that are comparable to values deduced from helioseismology, with both our models and apparently the sun possessing a contrast of order 30\\%.  Further, this is accomplished while imposing an upper thermal boundary condition that ensures a uniform emerging heat flux with latitude, as suggested in Elliott et al. (2000).  {\\sl Path 1} involving a decreasing Prandtl number is somewhat more effective in attaining large $\\Delta \\Omega$ as the solutions become turbulent than {\\sl Path 2} which has the Prandtl number fixed at 0.25 as both diffusivities are decreased.  This holds out hope that even more turbulent solutions will act likewise.   We have shown that the strong $\\Delta \\Omega$ results from the role of the Reynolds stresses in redistributing the angular momentum.  This transport is established by correlations in velocity components arising from convective structures that are tilted toward the rotation axis and depart from the local radial direction and away from the meridional plane.  These yield both $v_r v_{\\phi}$ and $v_{\\theta} v _{\\phi}$ correlations necessary for the Reynolds stress contributions to the radial and latitudinal angular momentum fluxes analyzed in Figure \\ref{fig11}. The fast downflow plumes have a dominant role in such Reynolds stresses, much as seen in local studies (Brummell et al. 1998).  Our simulations have attained a spatial resolution adequate to begin to attain coherent structures amidst the turbulence, which is believed to be a key in sustaining strong Reynolds stresses at higher turbulence levels.  This has the consequence that all our spherical shells possess fast prograde equatorial rotation relative to the reference frame. There are some contributions toward maintaining the differential rotation from the latitudinal heat transport inherent in our convection that serves to establish a warm pole (with a contrast of a few K) relative to the equator, with baroclinicity and a partial thermal--wind balance helping to yield equatorial acceleration.  The meridional circulations generally work to oppose such tendencies by redistributing angular momentum so as to try to spin up the poles.  Our simulations on {\\sl Paths 1} and {\\sl 2} confirm that strong differential rotation with fast equators has its primary origin in angular momentum transport associated with the Reynolds stresses.  Such prominent transports serve to resolve {\\sl Issue 2}.  Our next challenge is to satisfy {\\sl Issue 1} simultaneously with {\\sl Issue 2} in the more turbulent solutions, which may also lead to $\\Omega$ being more nearly constant on radial lines at mid to high latitudes.  Although our results for $\\Omega$ have made promising contacts with helioseismic deductions about the state of solar differential rotation in the bulk of the convection zone, there are also major issues that we have not yet tackled.  We must evaluate more advanced subgrid-scale terms in representing the unresolved turbulence within such simulations, especially in the near-surface regions.  Foremost are also questions of how does the presence of a region of penetration below the convection zone influence the angular momentum redistribution in the primary zone above, and does the tachocline of shear that is established near the interface with the deeper radiative interior modify properties within the convection zone itself.  We are keen to also investigate aspects of the rotational shear evident close to the solar surface.  This is just now becoming computationally feasible, and involves extending our computational domain upward and beginning to resolve supergranular motions there, as discussed in DeRosa \\& Toomre (2001) in preliminary studies with thin shells.  We are still at early stages with our simulations using ASH to study turbulent convection in spherical shells, yet it is comforting that the mean differential rotation profiles realized in some of the simulations are beginning to capture many of the dominant features for $\\Omega$ deduced from the helioseismic probing.   We thank Nicholas Brummell, Marc DeRosa, Julian Elliott, Peter Gilman, Mark Miesch and Jean-Paul Zahn for useful discussions and comments during the writing phase of this paper, and a referee for encouraging us to clarify the objectives and thrust of the presentation.  This work was partly supported by NASA through SEC Theory Program grant NAG5-8133 and by NSF through grant ATM-9731676.  Various phases of the simulations with ASH were carried out with NSF PACI support of the San Diego Supercomputer Center (SDSC), the National Center for Supercomputing Applications (NCSA), and the Pittsburgh Supercomputing Center (PSC).  Much of the analysis of the extensive data sets was carried out in the Laboratory for Computational Dynamics (LCD) within JILA.}"
    },
    "0206/astro-ph0206125_arXiv.txt": {
        "abstract": "Self-interacting dark matter  was proposed by Spergel $\\&$ Steinhardt (2000) to alleviate two conflicts between  Cold Dark Matter (CDM) models and observations. Firstly,  CDM N-body simulations predict dark matter halo density profiles that diverge at the centre in disagreement with the constant density cores observed in late-type dwarf and Low Surface Brightness (LSB) galaxies. Secondly,  N-body simulations predict an overabundance of subhalos in the Galactic halo. Using a simple semi-analytical argument we show that weakly self-interacting dark matter models, which can produce halo cores of the sizes observed in dark matter dominated galaxies, are unable to reconcile the number of satellites in the Galactic halo with the observed number of dwarf galaxies in the Local Group. ",
        "introduction": "Recent improvements in observational and numerical techniques have allowed a comparison between predictions of the CDM scenario and observational data on galactic scales. The results point out discrepancies between predictions and observations. High-resolution N-body simulations have shown that, on scales comparable to the Local Group, the predicted number of subhalos is at least a factor of ten higher than the observed number of dwarf galaxies (Klypin et al. 1999, Moore et al. 1999a). This disagreement, usually called the ``satellite question'', can be attributed to the high core densities of satellite dark halos found in cosmological models (Navarro, Frenk $\\&$ White 1997; hereafter NFW). These densities, combined with a small central velocity dispersion (Fukushige $\\&$ Makino 1997), tend to stabilize the satellites against tidal disruption on galactic scales. Another discrepancy emerges when comparing the density profiles of dark matter halos predicted by numerical simulations with observations of HI rotation curves in dwarf galaxies (Moore 1994; Flores $\\&$ Primack 1994; Burkert 1995). Whereas observations show linearly rising rotation curves out to radii greater than 1 $h^{-1}$ kpc, indicating that the dark matter has a constant density core (soft core), cosmological simulations predict dark halo density profiles with $\\rho \\propto r^{-1.5}$ in the central parts (Moore et al. 1999b; Fukushige $\\&$ Makino 2001). Other N-body simulations appear to converge to halo density profiles described by $\\rho \\propto r^{-1}$ (Power et al. 2002). These two conflicts, which might be related, the excess of dark satellites and the soft core question, arise because the CDM N-body simulations predict dark matter halos with high core densities.  Each conflict taken individually may not be sufficient to invalidate CDM on galactic scales. Results derived from observed density profiles of the inner regions in galaxies are controversial, due to beam smearing effects in HI rotation curves (van den Bosch $\\&$ Swaters 2001), even though high-resolution observations of H$\\alpha$ also show shallower core densities than those predicted by CDM numerical simulations (e.g. de Blok $\\&$ McGaugh $\\&$ Rubin 2001; Marchesini et al. 2002). Several authors have attempted to reconcile the number of observed Local Group dwarf galaxies with the number predicted by CDM theory through conservative solutions within the framework of the current theory. Early work by Kauffmann, White $\\&$ Guiderdoni (1993), using semi-analytic models of galaxy formation, found most subhalos lacked a luminous component. Energetic mechanisms which are more efficient in low mass systems, such as feedback from evolving stars and heating by an ionizing UV background, were proposed to explain a decoupling of luminous and dark components for low mass dwarfs (Efstathiou 1992, Bullock et al. 2000, Gelato $\\&$Sommer-Larsen 1999, Thacker $\\&$ Couchman 2000). Another solution, proposed by Klypin (1999), suggested an identification of the missing satellites seen in numerical simulations with observed compact high-velocity clouds (Blitz et al. 1999). This proposal may be premature, since it is still unclear  whether the high-velocity clouds are galactic or extragalactic in nature. Comparisons of the dark satellite halos in CDM dominated simulations to the distribution of observed neutral hydrogen high-velocity clouds and compact high-velocity clouds were made by Putman $\\&$ Moore (2001). Recently, Stoehr et al. (2002) and Hayashi et al. (2002) suggested that the Galactic satellites could be identified with the most massive subhalos of CDM simulations. This would tend to support scenarios in which baryons are lost preferentially from low-mass halos for yet unknown reasons.  Still, the disagreement between observations and predictions might indicate that a revision to the CDM scenario is required. Self-interacting dark matter was proposed by Spergel $\\&$ Steinhardt (2000) to overcome the satellite question and the soft core question. In this model, dark matter particles experience  weak, non-dissipative, collisions on scales of kpc to Mpc for typical galactic densities. These collisions thermalize the inner regions of the dark halos, producing a soft core. In addition, the excess of subhalos predicted by the CDM models would be reduced. This model has attracted great attention. Numerical simulations (e.g. Burkert 2000; Yoshida et al. 2000; Moore et al. 2000; Firmani, D'Onghia $\\&$ Chincarini 2001;  Dav\\'{e} et al. 2001; D'Onghia, Firmani $\\&$ Chincarini 2002) demonstrated that, in this scenario, soft cores would form naturally after a collisional timescale by energy transport into the cold inner regions. However, after the initial expansion, the cores of isolated halos would evolve towards the core-collapse stage, with final central densities higher than those predicted by NFW (Burkert 2000; Kochanek $\\&$ White 2000). Ostriker (2000) and Hennawi $\\&$ Ostriker (2002) pointed out that self-interacting dark matter in a very weak cross section regime in the centers of galaxies reproduces supermassive black hole masses and their observed correlation with the velocity dispersion of the host bulges. However, they point out a possible inconsistency of the collisional scenario; indeed, the model would lead to the exorbitant growth of supermassive black holes, which consequently  imposes a very strict upper limit on the collisional cross section. Other limits on the cross section values were derived from the fundamental plane relation for ellipticals in clusters (Gnedin $\\&$ Ostriker 2001).  Using an analytical approach, this Letter explores whether weakly self-interacting dark matter is likely to reconcile the apparent overabundance of subhalos with the small number of visible satellites in the Local Group. Two different disrupting processes are explored: collisions and tidal stripping. This work assumes $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "In a hierarchical universe, high-resolution N-body simulations of Standard CDM models predict an excess of subhalos with respect to the number of dwarf galaxies observed in the Local Group (Klypin et al. 1999, Moore et al. 1999a).  To solve this conflict between predictions and observations a successful theory should reduce  the abundance of substructures at all radii. However, if the `satellite question' remains a problem of CDM models, than our semi-analytical argument proves that self-interacting dark matter is unable to solve it.  If the value of the cross section is chosen such that it reproduces soft cores with sizes observed in late-type dwarf and LSB galaxies and  assumed to decrease with the halo velocity dispersion, then collisions between particles are  effective in disrupting subhalos only within  25 $h^{-1}$ kpc from the Milky Way centre. As a result, only a small percentage of all substructures are destroyed by collisions and the overabundance is only slightly reduced, leaving the problem unsolved at all radii. Tidal stripping is more efficient than collisions in destroying subhalos within 40 $h^{-1}$ kpc from the Galactic centre, alleviating the problem at small radii. However, discrepancy between predictions and observations persists.  In summary, finding a process that is able to decrease the halo central density does not appear  to be sufficient for reducing the excess of dark satellite halos, especially for substructures placed at large distances from the Milky Way centre. Thus, weak self-interaction, which was originally proposed to solve the soft core question in centres of  dark matter dominated galaxies and the overabundance of subhalos in the Local Groups is unable to solve both questions simultaneously."
    },
    "0206/astro-ph0206313_arXiv.txt": {
        "abstract": "We analyze the infrared (6-100\\,\\mic) spectral energy distribution of the blue compact dwarf and metal-poor (Z=\\zsolar/41) galaxy \\sbs. With the help of DUSTY \\citep{Ive99}, a program that solves the radiation transfer equations in a spherical environment, we evaluate that the infrared (IR) emission of \\sbs~is produced by an embedded super-star cluster (SSC) hidden under 10$^5$ \\msolar~of dust, causing 30 mag of visual extinction.  This implies that one cannot detect any stellar emission from the 2$\\times10^6$ \\msolar\\ stellar cluster even at near-infrared (NIR) wavelengths.  The derived grain size distribution departs markedly from the widely accepted size distribution inferred for dust in our galaxy (the so-called MRN distribution, \\citet{Mat77}), but resembles what is seen around AGNs, namely an absence of PAH and smaller grains, and grains that grow to larger sizes (around 1 \\micron).  The fact that a significant amount of dust is present in such a low-metallicity galaxy, hiding from UV and optical view most of the star formation activity in the galaxy, and that the dust size distribution cannot be reproduced by a standard galactic law, should be borne in mind when interpreting the spectrum of primeval galaxies. ",
        "introduction": "\\label{introduction} The question of how the energy radiated by a very young burst of star formation is redistributed in the electromagnetic spectrum by the neighboring ISM is one with far-reaching implications.  Indeed, as it is generally assumed that the formation of galaxies should be signalled by violent bursts of star formation (see e.g. the reviews by \\citet{Sil01} or \\citet{Ell98}, and references therein), the answer to this question can help defining the best observing strategy to study primeval galaxies.  For the most massive objects, it is generally assumed that star formation proceeds as observed in Ultra-Luminous InfraRed Galaxies (ULIRGs, see the review by \\citet{San96}).  In these systems, we know that most of the energy emerges in the infrared, and this would tend to invalidate any result derived from optical-UV surveys of the distant Universe.  Yet studies on the local starburst galaxy population appear to indicate a correlation between the total infrared luminosity and the extinction as measured by the slope of the UV continuum \\citep{Meu95}.  Such a correlation, along with the establishment of an effective attenuation curve \\citep{Cal94}, offers the hope to address the question of galaxy formation with optical-UV instruments, thus circumventing an important problem of most current infrared and submillimeter instruments: their lower spatial resolution that makes the identification of counterparts and subsequent determination of redshifts problematic.  However a number of relatively recent discoveries on the properties of starburst galaxies and ULIRGs cast some doubt on the potential of this UV-IR/Submm relation and on the physical meaning of an attenuation curve for getting at the intrinsic UV luminosity of a starburst galaxy. Recent high-spatial resolution MIR instruments have revealed the existence of very bright super star-clusters (clusters containing a few thousand O stars, hereafter SSCs) that are nearly or absolutely absent from visible images, e.g. the deeply buried SSC found in the Antennae galaxy \\citep{mvc98}. This object produces about 20\\% of the total MIR emission of the whole galaxy and was shown by \\citet{gil00} to be a very young ($\\sim$ 4\\,Myr) SSC containing 1.6$\\times10^{7}$\\,\\msolar\\ of stars embedded in an $A_{V}=10$ cloud of dust. This is no longer an isolated case: the Wolf-Rayet dwarf galaxy \\hen\\ is an even more impressive example of the buried SSC phenomenology. \\citet{Kob99} showed that \\hen\\ contains extremely compact radio sources whose spectrum is optically thick at 5\\,GHz, which are interpreted as ultra-dense \\hii\\ regions created by dust-embedded SSCs each with $\\sim$750\\,O7{\\sc V} stars. Gemini/OSCIR high resolution MIR observations by \\citet{Vac02} showed that the radio SSCs are exactly coincident with the MIR emitting regions observed previously by \\citet{Stl97}; the SSCs generate {\\em almost all} of the MIR luminosity of the galaxy, and there is {\\em no overlap} between the MIR emitting regions and those detected in the K band. This is also true with the L and M bands \\citep{Sau02} implying a very high optical depth along the line of sight toward the SSCs. Another case where the infrared emission arises from a dust-embedded SSC with no optical counterpart is the dwarf galaxy NGC\\,5253 \\citep{Tur00, Gor01}.  Recent observations have shown that dust is present even in the most metal-deficient objects in amounts large enough to affect our ability to observe the star-formation process, namely \\zw\\ and \\sbs. In \\zw, still the most metal-poor galaxy known at Z=\\zsolar/50, the analysis of the H$\\alpha$/H$\\beta$ ratio by \\citet{Can02} indicates patches of dust inside the \\hii\\ regions that lead to $A_{V}=0.5$\\,mag in some places.  \\sbs, at \\zsolar/41, for which we are presenting new data, has the highest star formation rate of the two.  \\citet{Thu97}, based on HST images, argue that this galaxy is probably undergoing its first burst of star formation (but see \\citet{Ost01}, and consider that the aim of this paper is {\\em not} to discuss whether or not \\sbs\\ undergoes its first burst of star formation, but rather to show that the {\\em current} burst properties can shed light on phenomena possibly occurring in primeval galaxies.  In other words, \\sbs\\ is considered in this work as a laboratory to study primeval galaxies, but not as a primeval galaxy itself).  In HST images, young stars appears concentrated in 6 SSCs, each of them not older than 25 Myr and all located within a region smaller than 526 pc\\footnote{With H$_o$ = 75 km s$^{-1}$ Mpc$^{-1}$.  At this distance, 1${\\arcsec}$ is 263 pc.}. In the NIR, the emission originates mostly from a region coincident with two of these SSCs (the ground-based NIR image does not allow to precisely attribute the emission to the HST-detected SSCs).  The NIR spectrum indicates stellar populations younger than 5\\,Myr \\citep{Van00}.  The picture gets more complex when the MIR properties are considered as well: the galaxy is very bright in the MIR and its global MIR spectrum is quite unusual (\\cite{Thu99}, hereafter Paper~I). First, it is lacking the Unidentified Infrared Bands (UIB) commonly attributed to Polycyclic Aromatic Hydrocarbons or PAHs, \\citep{Leg84, All85}.  This is generally indicative of dust exposed to a strong radiation field that either destroys the UIB carriers or swamps their emission in that of the very small grains.  Second it shows a marked silicate absorption band at 9.7\\,\\mic, very unusual at the galaxy scale and indicative of a large dust column density, unexpected in such a low metallicity galaxy.  This peculiar spectrum led to the hypothesis that the MIR emission originates from a dust-enshrouded SSC. Subsequent ground-based observations by \\citet{Dale01} showed the MIR emission to be almost a point source coincident with the NIR emitting region; Contrary to Paper~I, these authors argued against the buried SSC case for \\sbs.  Thus whether \\sbs\\ contains one or more deeply buried SSCs remains an open question (many different $A_V$ have been determined for the SSCs of \\sbs, ranging from $A_V \\sim 0.55$ based on the Balmer decrement \\citep{Izo97} to $A_V \\sim 20.0$ based on MIR spectroscopy, Paper~I), and is worth returning to.  In section~\\S\\ref{observations} we present new GEMINI/OSCIR and ISOPHOT observations used in conjunction with the ISOCAM data to reconstruct the infrared spectral energy distribution of the galaxy. In section~\\S\\ref{model} we define and justify our assumptions regarding the modelling of radiation transfer in \\sbs. Our results are presented in section~\\S\\ref{results}, and their implications are discussed in section~\\S\\ref{discussion}. ",
        "conclusions": "We have modelled the infrared SED of the blue compact dwarf galaxy \\sbs\\ with DUSTY, which solves consistently the radiation transfer in a spherical distribution of dust.  From this modelling, we deduce that \\sbs\\ harbors a deeply embedded super-star cluster, effectively hidden under about 30 mag of visual extinction.  The low-metallicity of the galactic gas did not preclude the formation of the 10$^5$ \\msolar\\ of dust necessary to completely hide from the optical view the SSC. With 2$\\times10^6$ \\msolar\\ of stars and an age of probably less than 5 Myr, the SSC has not been able yet to pierce through the cocoon of dust and gas from which it formed, but it had a profound effect on the dust size distribution: the hardness of the radiation destroyed the smallest dust particles and the PAH but shocks did not alter yet the larger size grains up to 1\\,\\micron.  The standard MRN distribution, normally observed in quiescent galactic environment, cannot reproduce the IR data we have at hand.  Instead, the dust in the SSC environment is reminiscent of what we observe around AGN, emphasizing the role of density and radiation hardness on the dust grain size distribution.  If dust-enshrouded SSCs are commonly associated to starbursting environment, the star-formation rate deduced by looking at the rest-frame optical or UV should be taken with caution.  Even with IR or MIR information care should be taken to use a correct radiation transfer treatment and one should use an extinction law suited to the radiation and gas density of the observed source."
    },
    "0206/astro-ph0206255_arXiv.txt": {
        "abstract": "The 40 ks \\emph{Chandra} ACIS-S observation of A1367 provides new insights into small-scale structures and point sources in this dynamically young cluster. Here we concentrate on small-scale extended structures. A ridge-like structure around the center (``the ridge'') is significant in the \\chandra\\ image. The ridge, with a projected length of $\\sim$ 8 arcmin (or 300 h$_{0.5}^{-1}$ kpc), is elongated from northwest (NW) to southeast (SE), as is the X-ray surface brightness distribution on much larger scales ($\\sim$ 2 h$_{0.5}^{-1}$ Mpc). The ridge is cooler than its western and southern surroundings while the differences from its eastern and northern surroundings are small. We also searched for small-scale structures with sizes $\\sim$ arcmin. Nine extended features, with sizes from $\\sim$ 0.5$'$ to 1.5$'$, were detected at significance levels above 4 $\\sigma$. Five of the nine features are located in the ridge and form local crests. The nine extended features can be divided into two types. Those associated with galaxies (NGC 3860B, NGC 3860 and UGC 6697) are significantly cooler than their surroundings (0.3 - 0.9 keV vs. 3 - 4.5 keV). The masses of their host galaxies are sufficient to bind the extended gas. These extended features are probably related to thermal halos or galactic superwinds of their host galaxies. The existence of these relatively cold halos imply that galaxy coronae can survive in cluster environment (e.g., Vikhlinin et al. 2001). Features of the second type are not apparently associated with galaxies. Their temperatures may not be significantly different from those of their surroundings. This class of extended features may be related to the ridge. We consider several possibilities for the ridge and the second type of extended features. The merging scenario is preferred. ",
        "introduction": "Clusters of galaxies exhibit a range of X-ray morphologies from regular and relaxed to irregular and dynamically young (Forman \\& Jones 1982). Irregular clusters with no X-ray dominant galaxy are generally considered dynamically young systems. Studies of such dynamically young systems, in combination with studies of more evolved clusters, are very important for understanding cluster evolution and its impact on the evolution of galaxies (and {\\em vice versa}). A1367, with its cool gas temperature (3-4 keV; e.g., Donnelly et al. 1998; D98 hereafter), high fraction of spiral galaxies, irregular X-ray distribution and low central galaxy density, is the prototype of such a dynamically young system (Bahcall 1977; Jones et al. 1979; Bechtold et al. 1983; B83 hereafter). It is a nearby (z=0.022) X-ray-bright cluster. Thus, detailed X-ray studies can be performed and the results will be very important for understanding dynamically young clusters.  Based on the \\einstein\\ HRI (EHRI hereafter) observations, B83 found fifteen 1$'$ scale emission features in A1367 with luminosities of several times 10$^{41}$ ergs s$^{-1}$ (0.5 - 4.5 keV). Most features seemed to be associated with galaxies and were attributed to hot galactic coronae (B83). However, Canizares, Fabbiano \\& Trinchieri (1987) noted that most early-type galaxies identified in A1367 by B83 have X-ray luminosities that exceed those of other galaxies with similar optical luminosities by nearly two orders of magnitude. Thus, they suggested that those regions found by B83 might be due to clumpiness in the ICM and should not be identified with galaxy halos. Grebenev et al. (1995; G95 hereafter) applied a wavelet analysis to both \\rosat\\ PSPC and EHRI data. They detected 8 more extended features mainly from the PSPC observation (in a larger field of view - FOV), though they only detected 8 of 15 extended features claimed by B83. Lazzati et al. (1998; L98 hereafter) performed their own wavelet analysis on the same PSPC data. Their results do not show many extended features ($\\sim$ 7) and they only detected 4 of 11 extended features found by G95. All three analyses found extended features that are not apparently associated with galaxies though their nature is still vague.  Here we present the analysis of the recent \\chandra\\ observation of A1367. With \\chandra's superior spatial resolution ($\\sim$ 0.5$''$ at the aimpoint) and better sensitivity than previous missions, we can better understand these extended features. In this paper, we concentrate on the small-scale structures. The analysis of point sources (including those corresponding to the member galaxies) will be presented in paper II (Sun \\& Murray 2002). The \\chandra\\ data reduction is described in $\\S$2; the detection of the extended features is in $\\S$3; $\\S$4 describes the \\chandra\\ temperature map; $\\S$5 discusses the central ridge and a surface brightness discontinuity at its southern end; $\\S$6 and $\\S$7 discuss extended features; and $\\S$8 is the discussion while $\\S$9 is the summary. Throughout this paper we assume H$_{0}$ = 50 km s$^{-1}$ Mpc$^{-1}$ and q$_{0}$ = 0.5. These cosmological parameters correspond to a linear scale of 0.62 kpc/arcsec at the cluster redshift. All luminosities, densities and masses scale as h$_{0.5}^{-2}$, h$_{0.5}^{3/2}$ and h$_{0.5}^{-3/2}$ respectively. ",
        "conclusions": "The \\chandra\\ observation of A1367 shows complicated structures around its center: a ridge-like structure and some extended features. For features not associated with galaxies, even in the absence of thermal evaporation losses, ram pressure stripping and Rayleigh-Taylor instabilities may destroy them within 10$^{7}$ - 10$^{8}$ yrs (Nittmann, Falle \\& Gaskell 1982; Sarazin 1986). We present several possibilities for the nature of these features.  \\vspace{0.5cm} \\noindent a) Background clusters\\\\  One may argue that some features are background clusters at sufficiently high redshift that the galaxies are too faint to show up in a shallow optical exposure. However, based on the log N - log S relation for clusters (e.g., Rosati et al. 1995; Kitayama, Sasaki \\& Suto 1998), we expect to see approximately 0.2 clusters at flux levels higher than 1.0$\\times$10$^{-14}$ ergs s$^{-1}$ cm$^{-2}$ (0.5 - 2 keV) in the 10$'\\times$10$'$ field (the flux of the faintest extended feature - C5 - is about 1.4$\\times$10$^{-14}$ ergs s$^{-1}$ cm$^{-2}$). Thus, the probability of seeing two (or three) clusters when 0.2 are expected is 1.6$\\times$10$^{-2}$ (or 1.1$\\times$10$^{-3}$). Moreover, five of six features (C1 - C5) seem to be located in the ridge and are aligned with each other, rather than distributed randomly as expected for background clusters. Thus, we conclude that the possibility for one or more features to be background clusters is very small.  \\vspace{0.5cm} \\noindent b) Features with no galaxies sustained by dark matter halos with little light\\\\  C1 - C6 can be supported by massive dark matter halos corresponding to very faint galaxies. The masses needed to sustain gas clouds from thermal diffusion can be estimated as the following. The assumed dark matter halos are required to have masses high enough that GMm$_{\\rm H}$/R $>$ 3/2 kT, where R and T are the sizes and the temperature of the feature respectively. Thus, M $>$ 6.7$\\times$10$^{11}$ (T/1 keV) (R/20 kpc). First we check the binding masses needed for C7 - C9. Since we know there are member galaxies associated with them, their total masses can be estimated by assuming M/L = 7(M/L)$_{\\odot}$. The results are presented in Table 3. It is clear that the total masses of the three galaxies are all larger than the masses needed to bind their corresponding extended features. We estimated the binding masses needed for C1 - C6. The results are 1 - 3 $\\times$10$^{12}$ M$_{\\odot}$, which is $\\sim$ 10 times larger than the total mass of any member galaxy close to C1 - C6 unless the nearby member galaxies have a large mass-to-light ratio of $\\sim$ 50. However, the alignment of C1 - C5 is not easy to explain and their similar temperatures to their surroundings imply that they are more likely to be inhomogeneities in the cluster medium.  \\vspace{0.5cm} \\noindent c) Recently stripped galaxy halos\\\\  Takeda, Nulsen \\& Fabian (1984) suggested some galaxyless extended features reported by B83 might be the blobs of the stripped gas. Based on the results in Nittmann et al. (1982), they estimated that such blobs can survive for about 10$^{8}$ years before being destroyed by Rayleigh-Taylor instabilities. Toniazzo \\& Schindler (2001) pointed out that these gas halos would be accelerated radially by the cluster potential gradient. Based on the expression they gave for the distance traveled by the gas blobs, we estimate that the blobs will travel about 50 - 150 kpc radially inward before being destroyed. Within 50 - 150 kpc of C1 - C6, there are not many known member galaxies (Fig. 2b). Moreover, it is known from the above discussion that it is difficult to explain each of C1 - C5 by gas stripping of only one galaxy. This scenario also has problem to explain the alignment of C1 - C5.  \\vspace{0.5cm} \\noindent d) The fragmented core of a merging component\\\\  As shown in Fig. 2, on the scale of about 2 h$_{0.5}^{-1}$ Mpc, the X-ray emission is largely flattened along the SE-NW direction and it can be divided into two parts spatially: one larger part at the SE near the peak of the cluster X-ray emission, and the other smaller part at the NW where there is a clear galaxy concentration centered at NGC 3842. Nevertheless, the galaxy distribution in the SE does not show any condensation near the X-ray center. D98 performed $\\beta$-model fits to these two components and constructed a temperature map based on \\asca\\ data. They suggested an ongoing merging along the SE-NW direction with two subclusters corresponding to the two parts mentioned. This \\chandra\\ observation reveals a ridge around the center elongated in the same direction as that of the flattened X-ray emission on a scale of about 2 h$_{0.5}^{-1}$ Mpc. Five of six extended features with no associated galaxy (C1 - C5) lie along the ridge and form five crests within it, while C6 is still close to the ridge. The ridge may have originally been the core of a subcluster. It is elongated along the merging direction because of the interaction with another subcluster. Instabilities (e.g., Kelvin Helmholtz instability) make it fragment into clumps we now observe (C1 - C5, possibly C6). The discontinuity found at the southern end of the ridge suggests motion of the ridge toward the SE. If this is true, the ridge could not be the core of the SE subcluster since its core is supposed to move to the NW in the interaction with the NW subcluster. Moreover, the shape of the ridge implies it has been distorted to a very large extent, by ram pressure stripping or tidal forces. This is possible for an infalling gas cloud, but not likely for the core of the larger SE subcluster at the very early stage of merging. It is difficult to relate the ridge with the NW subcluster since they are separated spatially (Fig. 2 and 7) and the NW subcluster peaks around a galaxy concentration centered at NGC 3842. Therefore, the ridge would have to be a new merging component, perhaps a compact precursor of the NW subcluster penetrating the SE subcluster. When it passes the SE subcluster, it is elongated and fragmented by ram pressure stripping or tidal force. Simulations show that at the early stage of cluster evolution, clusters may accrete smaller subclusters along filaments (summarized by Bertschinger 1998). Similar examples in X-ray and optical are presented by West (1998) for Coma and Durret et al. (1998) for A85. More observations in wider fields will be helpful to clarify the nature of the ridge."
    },
    "0206/astro-ph0206288_arXiv.txt": {
        "abstract": "{The temporal properties of a sample of 498 bright gamma-ray bursts (GRBs) with durations between 0.05 s and 674 s were analysed.  The large range in duration (T$_{90}$) is accompanied by a similarly large range in the median values of the pulse timing properties including rise time, fall time, FWHM and separation between the pulses.  Four timing diagrams relating these pulse properties to T$_{90}$ are presented and show the power law relationships between the median values of the 4 pulse timing properties and T$_{90}$, but also that the power laws depend in a consistent manner on the number of pulses per GRB. The timing diagrams are caused by the correlated properties of the pulses in the burst and can be explained by a combination of factors including the Doppler boost factor $\\Gamma$, a viewing effect caused by a jet and different progenitors.  GRBs with similar values of T$_{90}$ have a wide range in the number of pulses. GRBs with the large number of short and spectrally hard pulses may occur either from a homogeneous jet with a higher average value of $\\Gamma$ or close to the axis of an inhomogeneous jet with higher values of $\\Gamma$ near the rotation axis. The less luminous GRBs with fewer pulses may originate further from the axis of the inhomogeneous jet. The pulses in GRBs have six distinctive statistical properties including correlations between time intervals, correlations between pulse amplitudes, an anticorrelation between pulse amplitudes and time intervals, and a link to intermittency in GRS 1915+105. The timing diagrams and correlated pulses suggest that GRBs are powered by accretion processes signalling jets from the formation of black holes. ",
        "introduction": "The cosmological origin of gamma-ray bursts (GRBs) requires an extraordinary amount of energy to be emitted in gamma-rays \\citep{piran:1999}. The source of this energy is thought to be a cataclysmic event involving mergers of compact objects such as neutron stars or a neutron star and black hole \\citep{pacz:1991,ruffjan:1999} or the formation of a black hole in massive stars \\citep{macfad:1999,popham:1999}. GRB light curves are complex and irregular \\citep{fishman:1995} and a range of techniques have been developed to elucidate their structure \\citep{mhlm:1994,nnb:1996,lee:2000,belss:2000,quilligan:2002}. In this paper we present the results of the analysis of pulses in a large sample including long and short GRBs. The main results are presented in section 3 including four timing diagrams and our analysis of pulse properties. These results along with the unique properties of the pulses in GRBs are interpreted in section 4 and used to provide additional insights to the emission process and the central engine that support jets from accretion into newly formed black holes.  \\begin{figure*} \\resizebox{\\textwidth}{!}{\\rotatebox{-90}{\\includegraphics{Dj223_f1rgb.eps}}} \\label{fig:grbs_pulses} \\vspace{.5em} \\caption{The timing diagrams in GRBs. The median values obtained from the lognormal distributions for the pulse timing parameters (a) t$_{\\rm r}$, (b) t$_{\\rm f}$, (c) FWHM and (d) $\\Delta$T are plotted versus duration T$_{90}$ for GRBs in three categories i.e. 1 $\\leq$ N $\\leq$ 3 (blue), 4 $\\leq$ N $\\leq$ 12 (red) and N $>$ 12 (green). The crosses signify the range covered in T90 and \\(e^{\\mu \\pm \\sigma}\\) for the lognormal distribution which includes 16\\% to 84\\% of the pulse values for that T$_{90}$ bin.  The upper diagonal line is the limit where the pulse parameter is equal to T$_{90}$ and the lower lines by the limited time resolution i.e. 5 ms, 64 ms and 256 ms. In some GRBs with few pulses separated by a long $\\Delta$T the value of T$_{90}$ which goes from 5\\% to 95\\% of the counts may not include a small peak before or after the $\\Delta$T yielding a value of $\\Delta$T $>$ T$_{90}$. This effect is seen in (d).  Small corrections were applied to the pulse timing locations of the crosses  that took into account the truncation of the distributions by the limited time resolution. The lines are the best fits to the median values and the power law indices are listed in Table 1. The timing values were obtained for pulses that were isolated from neighbouring pulses at $\\geq$ 50\\% whereas all pulses above 5 $\\sigma$ were used for $\\Delta$T.} \\end{figure*} ",
        "conclusions": "The data in Fig. 1 shows that the median values of t$_{\\rm r}$, t$_{\\rm f}$, FWHM and $\\Delta$T depend strongly on T$_{90}$ and also on the number of pulses.  GRBs with the same durations can have a wide variation in the number of pulses and the pulse properties. GRBs with a small number of pulses generally have slow pulses that are further apart while bursts with larger number of pulses have faster and spectrally harder pulses that are closer together \\citep{norris:2001}. In both cases the same pattern is observed except that the median values of t$_{\\rm r}$, t$_{\\rm f}$, FWHM and $\\Delta$T scale with T$_{90}$. Bursts with only one pulse and correlated pulse properties would lie on a line of slope unity, while bursts with many pulses lie on parallel lines if the pulse properties and $\\Delta$T are correlated. The correlated nature of the pulse parameters and $\\Delta$T provide additional information on the central engine and the emission process.  \\subsection{Number of pulses, pulse properties and jets}  The separation of GRBs based on N (Fig. 1) suggests a kinematic origin because the median values of the pulse timing parameters and $\\Delta$T scale in the same way and by about the same amount. In the internal shock model, an effect of this type may occur in homogeneous jet models where the degree of collimation and range in $\\Gamma$ depends on the mass at the explosion \\citep{kobay:2001} or in an inhomogeneous jet model by a viewing effect caused by collimated emission with higher average values of $\\Gamma$ close to the rotation axis, where there is lower baryon pollution \\citep{reemes:1994,rees:1999,sal:2000}. The transparency radius of the fireball r$_{t}$ varies as $\\Gamma^{-\\frac{1}{2}}$.  The profiles of pulses in bursts from off-axis will be slowed by (1) the reduced value of $\\Gamma$ and (2) the longer time for faster shocks to catch slower ones before reaching r$_{t}$. There will be fewer shocks to collide outside r$_{t}$ and hence generate lower luminosity bursts with slower and softer pulses because the additional shock amalgamation produces a narrower range in $\\Gamma$. Indeed this process may be a key factor in controlling the spread in $\\Gamma$ in GRBs.  The luminosity of BATSE bursts is not determined but several factors imply that GRBs with faster and harder pulses are more luminous. The luminosity-variability and luminosity-lag correlations for GRBs with known z infer that the more variable and spectrally harder bursts are more luminous \\citep{nmb:2000,sal:2000,feniram:2001,schaefer:2001,ioka:2001,schmidt:2001}. These correlations can be explained in an inhomogeneous jet model where the more luminous and variable bursts with higher N occur close to the axis and the slower, softer and less luminous bursts from further off-axis \\citep{quilligan:2002}. In addition, there is the strong case for jets in GRBs based on the impressive range of afterglow studies \\citep{frail:2001}. The unexpected effect is that the jet has left its imprint on the timing profile of the burst.  \\subsection{Variation of pulse properties and $\\Delta$T with T$_{90}$}  The power laws in Fig. 1 connect the pulse data in the two sub-classes for 1 $\\leq$ N $\\leq$ 3 and 3 $\\leq$ N $\\leq$ 12. The median values of the pulse timing parameters and $\\Delta$T scale with T$_{90}$ by about a factor of 10$^{3}$. This result may indicate a direct connection between the two sub-classes of bursts or that they coincidentally lie on the same power laws in the absence of sufficient data to independently determine the slopes for the short bursts.  The popular progenitors of GRBs range from mergers of compact objects such as neutron stars and black holes to collapsars and hypernovae in massive stars. A neat feature of most progenitor models is that they provide a generic scenario based on the formation of a black hole with a massive debris torus that is rapidly accreted and energises the jet by neutrino transport and magnetohydrodynamic processes.  The variation of the pulse parameters with T$_{90}$ maybe caused by the same emission mechanisms and progenitors.  If GRBs have the same progenitors, the short GRBs must have very high values of $\\Gamma$ to produce the short durations \\citep{piran:1999}. The high value of $\\Gamma$ is close to the upper bound allowed by the internal shock model \\citep{sapi:1997,lazzati:1999} but well below what can in principle be produced in a very low baryon environment \\citep{meszar:1997}.  An alternative interpretation is that there are two types of progenitors with pulses that coincidentally lie on the same power laws. The progenitors that arise from massive stars may not be capable of producing short GRBs.  The short bursts may come from mergers of neutron stars or neutron stars and black holes \\citep{ruffjan:1999} where the time structure in the bursts reflects the interaction of thin relativistic shocks with $\\Gamma \\sim$ 100 and duration of $\\sim$ 50 ms.  In this two progenitor scenario the high values of $\\Gamma$ can be reduced by about an order of magnitude. In addition the central engine and the environment before the photosphere may combine to smooth the shocks in a way that is a function of T$_{90}$.  \\subsection{Correlations between pulses}  There are a number of new results on pulses, from short and long GRBs, that constrain the emission process.  These include (1) the distributions of values of t$_{\\rm r}$, t$_{\\rm f}$, FWHM, pulse amplitude (PA) and pulse area of isolated pulses are not random but have lognormal distributions \\citep{quilligan:2002,sheila:2001,np:2001}, (2) the values of $\\Delta$T are lognormally distributed with a power-law excess of long time intervals or a Pareto-Levy tail for GRBs with T$_{90} >$ 2 s \\citep{quilligan:2002,sheila:2001,np:2001}, (3) there is an anticorrelation between the values of PA and FWHM \\citep{ramfen:2000,quilligan:2002}, (4) there is a positive correlation between the values of the PA that extend over many pulses \\citep{quilligan:2002,sheila:2001}, (5) there is also a positive correlation between the values of $\\Delta$T \\citep{quilligan:2002,sheila:2001}, (6) the PA of isolated pulses in GRBs with T$_{90} >$ 2 s and 64 ms resolution were compared with $\\Delta$T$^{-}$ and $\\Delta$T$^{+}$ because of a strong correlation between the pulse width and $\\Delta$T$^{-}$ in the galactic superluminal source GRS 1915+105 where outbursts are powered by accretion into a black hole \\citep{belloni:1997} and a correlation between long quiescent periods in GRBs and subsequent intervals of emission \\citep{rm:2001}. The PAs were found to be anticorrelated with $\\Delta$T$^{-}$ and $\\Delta$T$^{+}$ with a slightly larger anticorrelation for $\\Delta$T$^{-}$.  Any additional correlations between the pulse properties and $\\Delta$T$^{-}$ would be expected to be much weaker than in GRS 1915+105 because a similar intermittent pattern in the outbursts from the central engine would have to persist after shock interactions generated the outbursts.  The combination of the six results make the pulses in GRBs quite unique.  In addition long pulses are asymmetric and reach their maximum earlier in higher energy bands while shorter pulses tend to be more symmetric with negligible time lags between energy channels \\citep{nnb:1996}.  The three results discussed above seem to favour an internal shock model powered by a hyper-accretion process into a newly formed black hole \\citep{ruffjan:1999,popham:1999,zhang:2001}. The accretion process has been modelled and is sensitive to the rate at which material piles up around the black hole and the accretion time scale of the particles. The accretion rate might provide the overall control on the process that eventually generates the correlated outbursts from the jet. The pulse amplitudes and time intervals can be coupled because higher accretion rates cause larger outbursts that are closer in time, while lower accretion rates produce smaller and slower events that are further apart. It appears that variation in the rate of accretion, the thickness of the relativistic shocks and the viewing angle of the jet may be key factors in accounting for the observed durations and pulse properties in GRBs. There are many uncertain factors that influence the process. These include the viscosity of the particles, the mass and angular momentum of the disk, the mass and spin of the black hole and the energy extraction and collimation process."
    },
    "0206/astro-ph0206241_arXiv.txt": {
        "abstract": "{In this paper we  present new observational evidence that supports the presence of an extra source of continuum emission in the binary system \\object{$\\beta$~Lyrae}. New VLA and IRAM observations, together with published data from the literature and ISO archive data, allow us to build the Spectral Energy Distribution of the binary between $5\\times 10^{9}$Hz and $5\\times 10^{15}$Hz. The radio-millimeter part of the spectrum is consistent with free-free emission from a symbiotic-like  wind associated with the primary component and ionized by the radiation field of the hidden companion. Furthermore, we also consider the possibility that the observed radio flux originates from collimated radio structures associated with the mass gaining component and its disk (Conical thermal jets). An extrapolation of this emission to the far-IR part of the spectrum indicates that in both cases the free-free contribution at these frequencies cannot explain the observations and that the observed infrared excess flux is due principally to the secondary component and its associated disk. ",
        "introduction": "Even if the eclipsing binary \\object{$\\beta$~Lyrae} is one of the most studied stellar systems,  its  enigmatic spectroscopic and photometric behaviour is not fully understood. The  current view is a non-degenerate, semi-detached interacting binary system in the phase of large-scale mass transfer between components. The mass losing component  is a  B6-B8p II, while the  unseen mass gaining component is probably a B0V star,  embedded in an optically thick accretion disk (Hubeny \\& Plavec \\cite{Hupla91}). The presence of a large plasma cloud, surrounding both components, has been inferred from optical and UV emission lines (Batten \\& Sahade \\cite{Batten}; Hack et al. \\cite{Hack}) as well as from the analysis of UV light curves (Kondo et al. \\cite{Kondo}). Recently  the presence of  jet-like structures, probably related to the accretion disk,  has been shown  by the studies of Harmanec et al. (\\cite{Harmanec96}) and confirmed by Hoffman et al. (\\cite{Hoff1998}). \\object{$\\beta$~Lyrae} is also a well known radio source. First detected, in the early seventies, by Wade \\& Hjellming (\\cite{Wade72}) at 2.7 and 8.1 GHz, it was then monitored, at two frequencies,  by Gibson (\\cite{Gibson75}). The source always exhibited a thermal-like spectrum but Wright \\& Barlow (\\cite{Wri1975}) pointed  out that the observed slope of the radio spectrum ($\\alpha$=0.96) was intermediate between the slope expected for a simple \\ion{H}{ii} region and from a stellar wind, implying that the physics underlying the radio emission of \\object{$\\beta$~Lyrae} is more complicated than assumed in either of these models.  Recent high resolution MERLIN observations of \\object{$\\beta$~Lyrae} at 4.9 $GHz$ (Umana et al. \\cite{Umana}) have revealed an extended radio nebula around the system, whose brightness temperature ($1.1 \\times 10^{4}$K) confirms the thermal origin of the radio emission. Such a nebula can be re-conducted to a massive wind associated with the cooler primary ionized by the hotter secondary if the two-winds model by Mazzali et al. (\\cite{Mazza92}) is adopted.  To assess a clear picture of the radio properties of \\object{$\\beta$~Lyrae}, an analysis of its spectral energy distribution (SED) from radio to infrared  appears to be necessary. In this paper we  present multi-frequency VLA  and IRAM observations that will be combined with data  from the literature in order to investigate the origin of the radio emission of \\object{$\\beta$~Lyrae}.  In the following, assuming the most common notation, we will indicate as primary the mass losing component, as it is the most luminous in the optical region. ",
        "conclusions": "\\object{$\\beta$~Lyrae} is one of the best studied stellar systems: still no reliable model of the system, able to reproduce the huge amount of observational data, is available.\\\\ Very recently, a new generation of models, that take into consideration the details of the physics of the accretions disk, have been presented (Linnell et al. \\cite{Linnell1998}; Bisikalo et al. \\cite{Bisikalo}). Even if they provide a good model for the light curves in the near and far-UV and in the visible, they still do not predict the shape of the infrared light curve, and, in particular, they do not reproduce the secondary minimum that becomes deeper than the primary minimum at $\\lambda \\geq 4.8 \\mu m$. This leads to the conclusion that another component, besides stars plus accretion disk, must contribute to the observed flux, and Linnell (\\cite{Linnell2000}) proposes that the extra source of continuum radiation may be Thomson scattering of radiation from the gainer.  In this paper we have presented new observational evidence that support the presence of an extra-component in the already quite complex binary system. This extra component, which is the  origin of  the observed thermal radio flux, can be due to an extended stellar wind ionized by the strong ultraviolet flux of the secondary component or related to collimated structures associated with the gainer (conical thermal jets).   We further evaluated the contribution of this component  in the far-IR to extract the spectral energy distribution of the stars plus accretion disk system. Our results, in both hypotheses,  show  a power-law distribution up to IRAS frequencies indicating an accretion disk with a non-uniform density distribution and whose size varies with wavelength. To quantify this, the collection of new, good quality infrared light curves would be highly valuable.  The stellar wind or thermal jets,  as appearing from the obtained spectrum and its modelling, is optically thick up to IRAS-ISO frequencies. Still, at those frequencies we measure a significant excess that we attribute to the disk since it is  in agreement with the near-IR data. We may ask what kind of effect  the radio nebula may have  on the observability of the inner disk. In both the cases, the size of the radio source depends on  frequency as $R_\\mathrm{radio} \\propto \\nu^{-0.7}$. Thus, scaling the size measured at 4.8 GHz with MERLIN (see Table~\\ref{nebula}) we obtain that the radio nebula reaches dimensions comparable to those of the entire binary system at frequencies of the order of $10^{13}$~Hz, close to the IRAS spectral region.  This provides important constraints to the radio source associated with the binary system,   whose morphology should allow one to observe also the disk. The other possibility, that the disk is surrounding the nebula, can be discarded, since a clear eclipse is observed in the infrared.  Finally, we would like to point out that even in the case of stellar wind, the fact that the radiation field necessary to ionize the wind is provided by the  secondary component, which is embedded in the thick accretion disk  (Linnell \\cite{Linnell2000}), suggests that the ionization would take place mainly in the polar regions. Thus, in both the cases,  it is quite probable that the nebula has a bipolar morphology that the MERLIN observations (Umana et al. \\cite{Umana}) were not able to resolve. This bipolar morphology, which  should have a wide opening angle as suggested by the observed spectral index (Rodriguez et al. \\cite{Rod1990}), is compatible with the possibility of observing the inner accretion disk."
    },
    "0206/astro-ph0206077_arXiv.txt": {
        "abstract": "The 23~GHz emission lines from the NH\\3 rotation inversion transitions are widely used to investigate the kinematics and physical conditions in dense molecular clouds.  The line profile is composed of hyperfine components which can be used to calculate the opacity of the gas \\citep{ho83}.  For intrinsic linewidths of a few \\kms, the 18 magnetic hyperfine components blend together to form a line profile composed of five quadrupole hyperfine lines.  If the intrinsic linewidth exceeds one half of the separation of these quadrupole hyperfine components ($\\sim5-10$~\\kms) these five lines blend together and the observed linewidths greatly overestimate the intrinsic linewidths.  If uncorrected, these artificially broad linewidths will lead to artificially high opacities.  We have observed this effect in our NH\\3 data from the central 10 pc of the Galaxy where uncorrected NH\\3(1,1) linewidths of $\\sim30$ \\kms ~exaggerate the intrinsic linewidths by more than a factor of two \\citep{gen87}.  Models of the effect of blending on the line profile enable us to solve for the intrinsic linewidth and opacity of NH\\3 using the observed linewidth and intensity of two NH\\3 rotation inversion transitions.  By using the observed linewidth instead of the entire line profile, our method may also be used to correct linewidths in historical data where detailed information on the shape of the line profile is no longer available. We present the result of the application of this method to our Galactic Center data.  We successfully recover the intrinsic linewidth ($\\langle\\dvi\\rangle\\approx15$~\\kms) and opacity of the gas. Clouds close to the nucleus in projected distance as well as those that are being impacted by Sgr A East show the highest intrinsic linewidths. The cores of the ``southern streamer'' \\citep{ho91,coi99,coi00} and the ``50 \\kms'' giant molecular cloud (GMC, \\citet{gus81}) have the highest opacities. ",
        "introduction": "} The metastable (J=K) rotation inversion transitions of NH\\3 are effective tracers of dense molecular material ($\\rm{N_H}\\approx10^4 - 10^5$~cm$^{-3}$) \\citep{ho83}.  The NH\\3 rotation inversion transitions are most often used to trace dense cores of giant molecular clouds (GMCs), but they are also observed in the ``streamers'' near the Galactic Center \\citep{ho91,coi99,coi00,mcg01} which will be our example data for this paper.  With frequencies near 23 GHz (1.3 cm), these transitions are easily observable with radio telescopes such as the Very Large Array\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (VLA) with only minimal interference by the atmosphere.  Compared to high frequency tracers such as the 3~mm transitions of HCN(J=1-0) and HCO$^+$(J=1-0) which often have $\\tau>1$, these NH\\3 transitions tend to be more optically thin, even in the case of low temperatures.  As a result, spectra of NH\\3 rotation inversion transitions are usually free of self-absorption effects.  The NH\\3 rotation inversion line is composed of 18 magnetic hyperfine lines which are grouped into five main features (a detailed description of the NH\\3 rotation inversion transitions can be found in \\citet{tow75} and \\citet{ho83}). The spacing of the magnetic hyperfine lines within each of the five components is too close to resolve except in cool and quiescent dense cores where the intrinsic linewidths of the gas can be less than 1~\\kms.  The effect of blending of the 18 magnetic hyperfine components on the observed linewidth is discussed in detail in \\citet{ho77} and \\citet{bar98}.  For linewidths of a few \\kms, the magnetic hyperfine components within each of the five features will blend together to form five electric quadrupole hyperfine components consisting of a main line and two symmetric pairs of satellite lines.  In this paper, we focus on large intrinsic linewidths and therefore will only concern ourselves with this ``five-peak'' profile.  These quadrupole satellite lines are located from 10 to 30 \\kms ~from the main line (see Table \\ref{table} for relative spacings and intensities).  Comparison of the intensity of a satellite line to the main line enables a direct calculation of the opacity of the gas making the NH\\3 rotation inversion transitions especially useful.  The derived opacity can then be used to calculate $N_H$, masses, and also combined with line ratios of different transitions to determine the rotational temperature of the gas \\citep{ho83}.  The small frequency separation of the five quadrupole hyperfine lines means that for large intrinsic linewidths, $\\dvi$, the satellite and main hyperfine components will be blended into a single line.  Figure \\ref{prof.fig} shows line profiles for NH\\3(1,1) for a range of intrinsic linewidths for $\\tm\\ll1$.  For $\\dvi\\ga 5$~\\kms ~the five components are blended into a single profile and it becomes difficult to measure the relative intensities needed for opacity estimation. The situation is aggravated by the fact that an increase in opacity will result in an increase in the intensity of the satellite hyperfine lines relative to the main line and the blended profile will appear to broaden even more.  As a result, an increase in either $\\tm$ or $\\dvi$ will appear as an increase in the observed linewidth, $\\dvo$, and it is difficult to interpret the data.  Exaggerated estimates of opacity as a result of mis-fitting intrinsically broad lines can greatly limit the physical parameters that can be determined from the data.  For example, if NH\\3(1,1) gas has $\\dvi>5$~\\kms, then an increase in linewidth may be mistaken as an increase in opacity.  The overestimation of the opacity propagates into overestimations of the column density and ultimately the calculated mass of the gas.  In the case of our Galactic Center data, the observed linewidths are high and the challenge is to determine the extent of the contribution from shocked material with large intrinsic linewidths as well as that from high opacity gas.  Shocks are important for identifying the locations of possible interactions between the streamers of molecular material that surround the Galactic Center.  A mistaken calculation of the mass of a cloud or incorrect identification of a shock could affect our general understanding of the environment in the central parsecs of the Galaxy.  It is therefore obvious that we need an improved method to obtain independent estimates of the opacity and intrinsic linewidth of the gas.  In this paper, we present a method to disentangle the effects of large intrinsic linewidths and high opacities by comparing the observed linewidths and intensities of two NH\\3 transitions.  Because the exact shape of the line profile need not be known, this method can be easily applied to historical data where the line profiles are no longer available.  Throughout the paper, we combine the theoretical discussion with a test case example using our NH\\3(1,1), (2,2), and (3,3) data from the Galactic Center.  In \\S \\ref{corr}, we show the correlation between $\\dvo$ and $\\tm$ in our Galactic Center data when the two quantities are solved for in the traditional manner. Section \\ref{dvodvi} then presents a calculation of the dependence of $\\dvo$ on $\\dvi$ for NH\\3(1,1), (2,2), and (3,3) for the special case where $\\tm\\ll1$.  This is followed by the incorporation of variations in opacity. The full method for simultaneously solving for the opacity and intrinsic linewidth is presented in \\S\\ref{method}.  The paper concludes by presenting the maps of the best estimate of $\\dvi$ and NH\\3(1,1) opacity for our Galactic Center data.  The map of intrinsic linewidths shows $\\langle\\dvi\\rangle=15$ \\kms ~for NH\\3(1,1) which is in agreement with observations of other tracers.  In addition, the corrected opacities show that all except the most massive features in the map are optically thin.  The opacity and intrinsic linewidths of the observed features are then discussed. ",
        "conclusions": "}  In gas with large intrinsic linewidths, the five quadrupole hyperfine components of the NH\\3 rotation inversion transitions will blend into a single-line profile.  The observed linewidth will greatly overestimate $\\dvi$ and bias the opacity to high values.  We model the effect of $\\dvi$ on observed NH\\3(1,1), (2,2), and (3,3) linewidths for a range of opacities.  A detailed method for recovery of the intrinsic linewidth and opacity using $\\dvo$ and the intensity of two NH\\3 rotation inversion transitions is then presented.  We focus on the case where only the observed linewidth of the profile is known which makes the method particularly useful for correcting historical data for which the exact line profiles are not readily available. Application of this method to our data from the central 10 pc of the Galaxy results in the first independent measures of the opacity and intrinsic linewidth of NH\\3(1,1) near the nucleus.  The calculated intrinsic linewidths when blending is accounted for are reduced by a factor of $\\sim2$ from the initial observed linewidths to $\\langle\\dvi\\rangle=15$~\\kms ~and now agree with the other molecular tracers in the region, including NH\\3(3,3).  Gas along the edge of Sgr A East or in close projected distance to the nucleus appears to have increased intrinsic linewidths.  All except the most massive clouds in our images show $\\ta<1$.  This confirms the observation that even the lowest energy transitions are less affected by self-absorption than HCN(1-0) and HCO$^+$(1-0)."
    },
    "0206/astro-ph0206307_arXiv.txt": {
        "abstract": "{% We present a two-layer accretion disc model developed to simultaneously fit optical long baseline visibilities and spectral energy distributions of T~Tauri accretion discs.  This model allows us to access easily the physical conditions in the disc as the mid-plane or the surface temperature.  Our model includes viscous heating, absorption of stellar irradiation, and thermalisation with the surrounding medium.  The disc is modelled with concentric cylinders for which the vertical radiation transfer is computed using two layers with vertically averaged temperatures: the outer layer is heated by the stellar irradiation and by the inner layer, and the inner layer by viscous dissipation and by the outer layer.  We investigate three prescriptions for the geometrical thickness of the disc: it is either proportional the scale height (model~1), given ad hoc (model~2), or zero (model~3).  We then derive the disc structure in the case of the $\\alpha$ and $\\beta$ viscosity prescriptions, as well as for various optical thickness regimes of the disc.  This analytical model allows us to disentangle regions where the mid-plane temperature and the effective temperature are dominated by accretion from regions dominated by reprocessing of stellar light.  In the case of $\\alpha$-prescription, we find that the structure of model~2 gives predictions very close to those of numerical simulations from previous authors.  From the disc structure, we derive the spectral energy distributions, images and interferometric visibilities.  We analyse the influence of the disc parameters on the resulting structure and on the observable outputs.  We apply our model to interpret consistently the spectral energy distributions and visibilities of SU~Aur and FU~Ori for which interferometric data are available, and that are not known to be part of a multiple system.  We were not able to derive a consistent fit for T~Tau North, which might come from caveats in the flux correction from its South component, but were able to separately derive fits for its spectrum and its visibilities.  We find that even a single interferometric measurement at one infrared wavelength can bring a very strong constraint on disc models. We predict that future massive interferometric observations of accretion discs will provide a breakthrough in the understanding of accretion disc physics. ",
        "introduction": "Since the initial models of viscous accretion discs by \\citet{Shakura73} and \\citet{LyndenBell74}, the physics of the close environment of T~Tauri stars (TTS) has been extensively studied in order to interpret their spectral energy distribution (SED).  For the sake of simplicity, models traditionally separated discs into two categories, sometimes called active discs, on one hand, in which viscous dissipation is predominant, and passive discs, on the other hand, for which irradiation by the central star is the main heating process.  Early models used quasi-Keplerian steady accretion discs, assumed to be geometrically thin for a wide range of accretion rate; they predicted a fixed slope for the infrared spectrum: $\\lambda F_\\lambda \\sim \\lambda^{-4/3}$. However many TTS present flatter SEDs, and disc flaring was among the first attempts to explore disc vertical structure as an explanation for such SED flattening \\citep{Adams87,Kenyon87}: in a flared disc, the surface of the remote parts is tilted toward the star and gets more stellar light than forecast by the standard model, resulting in a warmer disc further away from the star. Since then, several models have been proposed in order to explain both standard SEDs and flatter ones.  An analytical study of the radiative transfer in the vertical structure of discs was first carried out by \\citet{Hubeny90} for active discs, then by \\citet[hereafter paper~I]{paper1} for passive discs.  In the latter model, the topmost layers of the disc, illuminated by the star, are hotter than the disc photosphere, resulting in excess continuum and line emission.  Later on, \\citet{Chiang97,Chiang99} used a simplified two-layer passive disc model based on the same super-heating mechanism and derived SEDs, confirming conclusions of paper~I and producing results consistent with observations.  More recently, \\citet[hereafter paper~II]{paper2} generalised the analytical study of paper~I to discs heated by several processes, and used its formalism to derive the the vertical structure of active discs.  On the other hand, numerical integration of the equations of radiative transfer was carried out by various authors in order to derive the vertical structure of accretion discs with fewer \\latin{a priori} approximations. \\citet{Bell94,Bell97} developed an active disc model in order to explain FU~Orionis outbursts; \\citet{DAlessio98,DAlessio99} dealt with the more general case of a disc heated both by viscous heating and stellar irradiation.  From a general point of view, all these studies predict spectra consistent with observations, but they have rarely been checked consistently against the spatial information revealed by recent optical and infrared high angular resolution imaging.  Recently, the advent of optical interferometry has set newer constraints on disc models.  \\citet{Malbet98,Akeson00,Akeson02,Malbet01p} obtained the first visibility measurements of TTS and FU~Orionis stars. However, they failed to consistently fit both SEDs and visibilities with a standard disc model: most of the time, the disc parameters derived from the SED data are in disagreement with those derived from visibility data.  No other attempt to compare self-consistent disc models and interferometric measurements has been carried out so far for low mass pre-main sequence stars.  In this paper, we tackle the issue of analytically describing a disc in presence of the two main heating processes, viscous heating and stellar irradiation, the latter requiring a correct description of the flaring.  This model suits the TTS and FU~Ori-type stars for which viscous heating cannot be ignored; for more massive stars (Herbig Ae/Be) it has been shown \\citep{Dullemond01} that viscosity, as a heating mechanism, can be ignored, so that our model is not relevant.  In Sect.~\\ref{sec:model}, we present a two-layer version of the model developed in paper~I and carry out an analytical determination of the structure of the disc. We derive a set of equations giving the mid-plane temperature and the flaring index, from which the whole structure can be determined.  In Sect.~\\ref{sec:num}, we briefly present the numerical approach.  In Sect.~\\ref{sec:ana}, we compare the results of the model with other models and analyse the influence of some disc parameters on the observables.  In Sect.~\\ref{sec:res}, we apply our model to the few low-mass young stellar objects (YSOs) observed in optical interferometry. ",
        "conclusions": "\\label{sec:ccl}  We have developed an analytical model for T~Tauri accretion discs based on a two-layer approximation, and including the main heating processes: viscous dissipation, reprocessing, and thermalization with the surrounding medium.  The outer layer is directly heated by visible stellar light and the thermal flux from the inner layer, whereas the inner layer is heated by viscous dissipation and light reprocessed by the outer layer.  The strength of the model is an analytical prediction of the mid-plane temperature and, yet with less accuracy, of the flaring of the disc; it is a suitable tool in grasping the physical conditions taking place in these discs: it allows to predict easily the relative importance of heating processes, the contribution of scattering,  the influence of the viscosity prescription, etc.  Despite of simplifications made in order to keep the model analytical, it compares well with other disc models available in the literature.  Its predictions in terms of mid-plane temperature, density scale height or disc mass are consistent with those of numerical models by \\citet{DAlessio98}, or by \\citet{Bell97} in the inner regions where this model is valid.  For the first time, we are able to consistently fit the infrared spectra and the optical visibility of two young stars, SU~Aur and FU~Ori.  A third object has been considered (T~Tau North), but we could not find a set of parameters that would fit simultaneously the SED and the visibility.  This might result from the peculiar structure of this object which is a triple system, and demonstrates that even a single interferometric measurement at one infrared wavelength can be a very strong constraint to interpret disc models.  We therefore expect a breakthrough in disc physics understanding when new generation interferometers, like the VLTI or the KI, are able to observe hundreds of young stellar objects.  Providing an analytical description and a fast computation, while including essential physical phenomena taking place in discs, our model will be a useful tool to interpret these forthcoming observations.  Concerning the influence of the viscosity on the model output, our results show that current instruments cannot significantly probe the influence of the viscosity law, because it has no direct impact on the flux emerging from optically thick regions.  The dependency on viscosity observed in our model remains small and results from the variation of the disc geometrical thickness when the amount of material changes.  It is only at larger wavelengths, where the central regions are optically thin, that the mid-plane temperature can be probed, thus indirectly the viscosity.  The future Atacama Large Millimetre Array (ALMA) should be able to measure the column density and the mid-plane temperature of discs within 10 AUs from the star, and probe the influence of the viscosity.  In forthcoming work, we will address the following points. \\begin{itemize} \\item Including direct heating of the inner rim of the disc by the star will allow us to make predictions for more massive stars (Herbig Ae/Be systems) as shown by \\citet{Dullemond01}. \\item The treatment of self-shadowing can also be improved: our model~1 assumes that the regions where the disc flares inwards are simply not illuminated at all, while vertical structure models tend to prove that there is still enough material at large heights to catch stellar radiation.  Model~2 is neither satisfactory because the flaring index is taken \\latin{ad hoc}, though it gives quite faithful predictions. \\item A precise determination of the chemistry species in the outer parts of the disc implies to model the temperature profile in the outer layer instead of assuming an isothermal one.  \\citet{Aikawa02} showed that the model by \\citet{Chiang97} cannot reproduce observed element abundances in the outer part of the disc whereas the vertical structure by \\citet{DAlessio98} can. \\end{itemize}"
    },
    "0206/astro-ph0206131_arXiv.txt": {
        "abstract": "This pedagogical review addresses several issues related to statistical description of gravitating systems  in both static and expanding backgrounds, focusing on the latter. After briefly reviewing the results for the static background, I describe the key issues and recent progress in the context of gravitational clustering\\index{gravitational clustering}  of collision-less particles in an expanding universe\\index{expanding universe}. The questions addressed include: (a) How does the power injected into the system at a given wave number  spread to smaller and larger scales? (b) How does the power spectrum of density fluctuations\\index{density fluctuations} behave asymptotically at late times? (c) What are the universal characteristics  of gravitational clustering that are independent of the initial conditions and manifest at the late time evolution of the system? The review is intended for non cosmologists and will be of interest to people working in fluid mechanics, non linear dynamics and condensed matter physics. \\index{abstract} ",
        "introduction": "Introduction }  The statistical mechanics of systems dominated by gravity is of interest both from the theoretical and ``practical''  perspectives. Theoretically this field has close connections with areas of condensed matter physics, fluid mechanics, re-normalization group etc. and poses an incredible challenge as regards the basic formulation. From the practical point of view, the ideas find application in many different areas of astrophysics and cosmology, especially in the study of globular clusters, galaxies and gravitational clustering in the expanding universe. (For a general review of statistical mechanics of gravitating systems, see \\cite{tppr};  textbook description of the subject is available in \\cite{textone}, \\cite{texttwo}. Review of gravitational clustering in expanding background is available in \\cite{tpiran} and in several textbooks in cosmology \\cite{cosmotext}. There have been many attempts to understand these phenomena by different groups; see \\cite{chavanis}, \\cite{sanchez}, \\cite{vala}, \\cite{fola}, \\cite{botta}, \\cite{roman}  and the references cited therein.) Given the diversity of the subject, it will be useful to begin with a broad overview and a  description of the issues which will be addressed in this review.  I will  concentrate mostly on the problem of gravitational clustering in the expanding universe which is one of the most active research areas in cosmology. However, to place this problem in context, it is necessary to discuss statistical mechanics of isolated gravitating systems (without any cosmological expansion) in some detail. In  Part I of this review, I cover this aspect highlighting the important features but referring the reader to existing previous literature for details. Part II presents a  more detailed description of gravitational clustering in the context of cosmology. The rest of the introduction will be devoted to a  summary of different issues which will be expanded upon in the later sections.  Let me begin with the issues which arise in the study of isolated gravitating systems (like, say, a cluster of stars) treated as a collection of  structure-less point particles. In Newtonian theory, the gravitational force can be described as a gradient of a scalar potential and the evolution of a set of particles under the action of gravitational forces can be described the equations \\begin{equation} \\ddot{\\bf x}_i = - \\nabla \\phi ({\\bf x}_i, t); \\quad \\nabla^2 \\phi = 4 \\pi G \\sum_i m_i \\delta_D ({\\bf x} -{\\bf x}_i) \\end{equation} where  ${\\bf x}_i$ is the position of the $i-$th particle, $m_i$ is its mass.  For sufficiently large number of particles, it is useful to investigate whether some kind of statistical description  of such a system is possible. Such a description, however,  is complicated by  the long range, \\index{long range} unscreened, nature of gravitational force.  The force acting on any given particle receives contribution even from particles which are  far away. If a self gravitating system is divided into two parts, the total energy of the system cannot be expressed as the sum of the gravitational energy of the components. Many of the conventional results in statistical physics are based on the extensivity of the energy \\index{extensivity of the energy} which is clearly not valid for gravitating systems. To make progress, we have to use different techniques which are appropriate for each situation.  To construct the statistical description of such a system, one should begin with the construction of the micro-canonical ensemble describing such a system. If the Hamiltonian of the system is $H(p_i, q_i)$ then the volume $g(E)$ of the constant energy surface  $H(p_i, q_i) =E$ will be of primary importance in the  micro-canonical ensemble\\index{micro-canonical ensemble}. The logarithm of this function will give the entropy $S(E) = \\ln g(E) $ and the temperature of the system will be $T(E)\\equiv \\beta(E)^{-1} = (\\pder{S}{E})^{-1}$.  Systems for which a description based on canonical ensemble is possible, the Laplace transform of $g(E)$ with respect to a variable $\\beta$ will give the partition function $Z(\\beta)$. It is, however, trivial to show that gravitating systems of interest in astrophysics cannot be described by a canonical ensemble \\cite{tppr}, \\cite{dlbone}, \\cite{dlbtwo}. Virial theorem holds for such systems  and we have $(2K+U) =0$ where $K$ and $U$ are the total kinetic and potential energies of the system. This leads to $E=K+U= -K$; since the temperature of the system is proportional to the total kinetic energy, the  specific heat\\index{specific heat} will be negative: $C_V \\equiv (\\pder{E}{T})_V \\propto (\\pder{E}{K}) < 0$.   On the other hand, the specific heat of any system described by a canonical ensemble $C_V = \\beta^2 <(\\Delta E)^2>$ will be positive definite. Thus one cannot describe self gravitating systems of the kind we are interested in by  canonical ensemble\\index{canonical ensemble}.  One may attempt to find the equilibrium configuration for self gravitating systems by maximizing the entropy $S(E)$ or the phase volume $g(E)$. It is again easy to show that no global maximum for the entropy exists for classical point particles interacting via Newtonian gravity. To prove this, we only need to construct a configuration with   arbitrarily high entropy which can be achieved as follows: Consider a system of $N$ particles initially occupying a region of finite volume in phase space and total energy $E$. We now move  a small number of these particles (in fact, a pair of them, say, particles 1 and 2  will do) arbitrarily close to each other. The potential energy of interaction of  these two particles, $-Gm_1m_2/r_{12}$, will become arbitrarily high as $r_{12}\\to 0$. Transferring some of this energy to the rest of the particles, we can increase their kinetic energy without limit. This will clearly increase the phase volume occupied by the system without bound. This argument can be made more formal by dividing the original system into a small, compact core and a large diffuse halo and allowing the core to collapse and transfer the energy to   the halo.   The absence of the global maximum for entropy --- as argued above --- depends on the idealization that there is no short distance cut-off in the interaction of the particles, so that we could take the limit $r_{12}\\to 0$. If we assume, instead, that each particle has a minimum radius $a$, then the typical lower bound to the gravitational potential energy contributed by a pair of particles will be $-Gm_1m_2/2a$. This will put an upper bound on the amount of energy that can be made available to the rest of the system.  We have also assumed that part of the system can expand without limit --- in the sense that any particle with sufficiently large energy can move to arbitrarily large distances. In real life, no system is completely isolated and eventually one has to assume that the meandering particle is better treated as a member of another system. One way of obtaining a truly isolated system is to confine the system  inside a spherical region of radius $R$ with, say, reflecting wall.  The two cut-offs $a$ and $R$ will make the upper bound on the entropy finite, but even with  the two cut-offs, the primary nature of gravitational instability cannot be avoided. The basic phenomenon described  above  (namely, the formation of a  compact core\\index{compact core} and a  diffuse halo\\index{diffuse halo}) will still occur since this is the direction of increasing entropy. Particles in the hot diffuse component will permeate the entire spherical cavity, bouncing off the walls and having a kinetic energy which is significantly larger than the potential energy. The compact core will exist as a gravitationally bound system with very little kinetic energy. A more formal way of understanding this phenomena is based on  the virial theorem for a system with a short distance cut-off confined to a sphere of volume $V$. In this case, the virial theorem will read as \\cite{textone} \\begin{equation} 2T + U = 3PV + \\Phi \\label{modvirial} \\end{equation} where $P$ is the pressure on the walls and $\\Phi$ is the correction to the potential energy arising from the short distance cut-off. This equation can be satisfied in essentially three different ways. If $T$ and $U$ are significantly higher than $3PV$, then we have $2T + U \\approx 0$ which describes a self gravitating systems in standard virial equilibrium but not in the state of maximum entropy. If $T \\gg U$ and $3PV\\gg \\Phi$, one can have $2T \\approx 3PV$ which describes an ideal gas with no potential energy confined to a container of volume $V$; this will describe the hot diffuse component at late times. If  $T \\ll U$ and $3PV\\ll \\Phi$, then one can have $U\\approx \\Phi$ describing the compact potential energy dominated core at late times. In general, the evolution of the system will lead to the production of the core and the halo and each component will satisfy the virial theorem in the form (\\ref{modvirial}). Such an asymptotic state with two distinct phases  \\cite{aaron}  is quite different from what would have been expected for systems with only short range interaction. Considering its importance, I shall briefly describe in section \\ref{sec:phases} a toy model which captures the essential physics of the above system in an exactly solvable context.  The above discussion focussed on the existence of global maximum to the entropy and we proved that it does not exist in the absence of two cut-offs. It is, however, possible to have {\\it local} extrema of entropy which are not global maxima. Intuitively, one would have expected the distribution of matter in  the configuration which is a local extrema of entropy to be described by a Boltzmann distribution, with the density given by $\\rho ({\\bf x}) \\propto \\exp[-\\beta \\phi({\\bf x})]$ where $\\phi$ is the gravitational potential related to $\\rho$ by Poisson equation. This is indeed true and a formal proof will be given in section \\ref{mfgrav}.  This configuration is usually called the  isothermal sphere\\index{isothermal sphere} (because it can be shown that, among all solutions to this equation, the one with spherical symmetry maximizes the entropy) and since it is a local maximum of entropy, it deserves careful study. I will describe briefly some of the interesting features of the isothermal spheres in section \\ref{isosph}   and  this configuration will play a dominant role throughout our review. The second (functional)  derivative of the entropy with respect to the configuration variables will determine whether the local extremum of entropy is a local maximum or a saddle point \\cite{antonov}, \\cite{tpapjs} and some of these  results are described at the end of section \\ref{isosph}. The relevance of the long range of gravity in all the above phenomena can be understood by studying model systems with an attractive potential varying as $r^{-\\alpha}$ with different values for $\\alpha$. Such studies confirm the results and interpretation given above; (see \\cite{ispo} and references cited therein).  Let us now consider the situation in the context of an expanding background described in Part II which is the main theme of the review. There is considerable amount of observational evidence to suggest that one of the dominant energy densities in the universe is contributed by self gravitating point particles. The smooth average energy density of these particles drive the expansion of the universe while any small deviation from the homogeneous energy density will cluster gravitationally. One of the central problems in cosmology is to describe the non linear phases of this gravitational clustering starting from a initial spectrum of density fluctuations. It is often enough (and necessary) to use a statistical description and  relate different statistical indicators (like the  power spectra\\index{power spectra}, $n$th order  correlation functions\\index{correlation functions} ....) of the resulting density distribution to the statistical parameters (usually the power spectrum) of the initial distribution. The relevant scales at which gravitational clustering is non linear are less than about 10 Mpc (where 1 Mpc = $3\\times 10^{24}$ cm is the typical separation between galaxies in the universe) while the expansion of the universe has a characteristic scale of about few thousand Mpc. Hence, non linear gravitational clustering in an expanding universe can be adequately described by Newtonian gravity provided the rescaling of lengths due to the background expansion is taken into account. This is easily done by introducing a  {\\it proper} coordinate\\index{proper coordinate} for the $i-$th particle ${\\bf r}_i$, related to the  {\\it comoving} coordinate\\index{comoving coordinate}  ${\\bf x}_i$, by ${\\bf r}_i = a(t) {\\bf x}_i$ with $a(t)$ describing the stretching of length scales due to cosmic expansion. The Newtonian dynamics works with the proper coordinates ${\\bf r}_i$ which can be translated to the behaviour of the comoving coordinate ${\\bf x}_i$ by this rescaling. (Some basic results in cosmology are summarized in Appendix A.)  As to be expected, cosmological expansion  completely changes the nature of the problem because of several new factors which come in: (a) The problem has now become time dependent and it will be pointless to look for equilibrium solutions in the conventional sense of the word. (b) On the other hand, the expansion of the universe has a civilizing influence on the particles and  acts counter to the tendency of gravity to make systems unstable. (c) In any small local region of the universe, one would assume that the conclusions describing a finite gravitating system will still hold true approximately. In that case, particles in any small sub region will be driven towards configurations of local extrema of entropy (say, isothermal spheres) and towards global maxima of entropy (say, core-halo configurations).  An extra feature comes into play as regards the expanding halo from any sub region. The expansion of the universe acts as a damping term in the equations of motion and drains the particles of their kinetic energy --- which is essentially the lowering of temperature of any system participating in cosmic expansion.  This, in turn, helps gravitational clustering since the potential wells of nearby sub regions can capture particles in the expanding halo of one region when the kinetic energy of the expanding halo  has been sufficiently reduced.  The actual behaviour of the system will, of course, depend on the form of $a(t)$. However, for understanding the nature of clustering, one can take $a(t) \\propto t^{2/3}$ which describes a matter dominated universe with critical density (see Appendix A). Such a power law has the advantage that there is no intrinsic scale in the problem. Since Newtonian gravitational force is also scale free, one would expect some scaling relations to exist in the pattern of gravitational clustering. Incredibly enough, converting this intuitive idea into a concrete mathematical statement turns out to be non trivial and difficult. I shall discuss several attempts to give concrete shape to this idea in sections \\ref{renormgrav}, \\ref{nlscales} and \\ref{univgravcl} but there is definite scope for further work in this direction.   To make any progress we need a theoretical formulation which will  relate statistical indicators in the non linear regime of clustering to the initial conditions. In particular, we need a robust prescription which will allow us to obtain the two-point correlation function and the nonlinear  power spectrum from the initial power spectrum. Fortunately, this problem has been solved to a large extent and hence one can use this as a basis for attacking several other key questions. There are four key theoretical questions which are of considerable interest in this area that I will focus on:   \\begin{itemize} \\item If the initial power spectrum is sharply peaked in a narrow band of wavelengths, how does the evolution transfer the power to other scales? This is, in some sense, analogous to determining the  Green function for the gravitational clustering\\index{Green function for the gravitational clustering} except that superposition will not work in the non linear context. \\item What is the asymptotic nature of evolution for the self gravitating system in an expanding background? In particular, how can one connect up the local behaviour of gravitating systems to the overall evolution of clustering in the universe? (If we assume that the isothermal spheres play an important role in the local description of gravitating system, we would expect a strong trace of it to survive even in the context of cosmological clustering. This is indeed true as I shall show in sections \\ref{nlscales}, \\ref{cipt} and \\ref{univgravcl} but only in the asymptotic limit, under certain assumptions.) \\item Does the gravitational clustering at late stages wipe out the memory of initial conditions or does the late stage evolution depend on the initial power spectrum of fluctuations? \\item Do the  virialized structures\\index{virialized structures} formed in an expanding universe due to gravitational clustering have any invariant properties? Can their structure be understood from first principles? \\end{itemize}   All the above questions are, in some sense, open and thus are good research problems. I will highlight the progress which has been made and give references to original literature for more detailed discussion.  \\vskip 1cm  \\noindent {\\Large{\\bf Part I: Gravitational clustering in static backgrounds  }}  \\vskip 0.2cm ",
        "conclusions": ""
    },
    "0206/astro-ph0206182_arXiv.txt": {
        "abstract": "{Recently, cosmic shear, the weak lensing effect by the inhomogeneous matter distribution in the Universe, has not only been detected by several groups, but the observational results have been used to derive constraints on cosmological parameters. For this purpose, several cosmic shear statistics have been employed. As shown recently, all two-point statistical measures can be expressed in terms of the two-point correlation functions of the shear, which thus represents the basic quantity; also, from a practical point-of-view, the two-point correlation functions are easiest to obtain from observational data which typically have complicated geometry. We derive in this paper expressions for the covariance matrix of the cosmic shear two-point correlation functions which are readily applied to any survey geometry. Furthermore, we consider the more special case of a simple survey geometry which allows us to obtain approximations for the covariance matrix in terms of integrals which are readily evaluated numerically. These results are then used to study the covariance of the aperture mass dispersion which has been employed earlier in quantitative cosmic shear analyses. We show that the aperture mass dispersion, measured at two different angular scales, quickly decorrelates with the ratio of the scales. Inverting the relation between the shear two-point correlation functions and the power spectrum of the underlying projected matter distribution, we construct estimators for the power spectrum and for the band powers, and show that they yields accurate approximations; in particular, the correlation between band powers at different wave numbers is quite weak. The covariance matrix of the shear correlation function is then used to investigate the expected accuracy of cosmological parameter estimates from cosmic shear surveys. Depending on the use of prior information, e.g. from CMB measurements, cosmic shear can yield very accurate determinations of several cosmological parameters, in particular the normalization $\\sigma_8$ of the power spectrum of the matter distribution, the matter density parameter $\\Omega_{\\rm m}$, and the shape parameter $\\Gamma$. ",
        "introduction": "Cosmic shear, the distortion of the images of distant galaxies by the tidal gravitational field of intervening matter inhomogeneities, offers a direct way of probing the statistical properties of the large-scale (dark) matter distribution in the Universe, without making any assumption on the relation between dark and luminous matter (e.g., Blandford et al. 1991, Miralda-Escude 1991, Kaiser 1992, 1998, Jain \\& Seljak 1997, Bernardeau et al. 1997, Schneider et al. 1998, hereafter SvWJK, van Waerbeke et al. 1999; Bartelmann \\& Schneider 1999; Jain et al. 2000, White \\& Hu 2000; see Mellier 1999 and Bartelmann \\& Schneider 2001 for recent reviews). The first detections of cosmic shear on wide-field imaging data (Bacon et al. 2000, 2002; Kaiser et al. 2000; van Waerbeke et al. 2000, 2001, 2002; Wittman et al. 2000, Maoli et al. 2001; Rhodes et al. 2001; H\\\"ammerle et al. 2002; Hoekstra et al.\\ 2002; Refregier et al.\\ 2002) has demonstrated the feasibility of this new window of observational cosmology, and yielded already the first constraints on cosmological parameters, most noticibly the normalization $\\sigma_8$ of the dark matter power spectrum, but also on the matter density parameter $\\Omega_{\\rm m}$ (van Waerbeke et al.\\ 2002; Hoekstra et al.\\ 2002).  Most analytical work on cosmic shear has been done on two-point statistical measures of the distortion field, such as the shear correlation functions, the shear variance in an apertures, or the aperture mass (see Sect.\\ts 2 for a definition of these quantitities). Although higher-order statistical measures, such as the skewness of the shear (Bernardeau et al.\\ 1997), are likely to yield additional, if not even superior constraints on cosmological parameters, their theoretical predictions are more uncertain at present. Recently, an estimator for the skewness of the shear was developed (Bernardeau et al.\\ 2002a), and applied to wide-field survey data (Bernardeau et al\\ 2002b), yielding a significant detection.  In this paper we consider second-order statistical measures only. All of them can be derived in terms of the correlation functions, as shown in, e.g., Crittenden et al.\\ (2002, hereafter C02) and Schneider et al.\\ (2002, hereafter SvWM), and since the measurement of the correlation functions is easier in practice than the other two-point statistics (e.g., gaps in the data are easily dealt with), we consider the correlation functions as the fundamental observables from a cosmic shear survey. In order to use them for determining cosmological parameters, it is important to know a practical and unbiased estimator for them, and to determine the covariance of this estimator. Two effects enter this covariance: a random part, which is due to the intrinsic ellipticity of the galaxies from which the shear is measured, together with measurement noise, and sampling (or cosmic) variance. The first of these effects is expected to dominate on small angular scales, whereas the second determines the covariance for large separations. The covariance will depend not only on the total survey area, but also on the survey geometry. As has been pointed out by Kaiser (1998), in order to decrease the sampling variance on large scales, it may be favourable to choose a survey geometry that samples those scales sparsely. In order to design an optimized survey geometry, the covariance as a function of survey geometry needs to be calculated.  Here, we calculate the covariance matrices for the shear correlation functions binned in angular separation. In Sect.\\ts 2, we introduce our notation and briefly summarize the two-point cosmic shear measures and their interrelations. Unbiased estimators of the two basic correlation functions are derived in Sect.\\ts 3, togther with the corresponding unbiased estimators of the aperture mass and the shear dispersion. The covariance matrices of these correlation function estimators are then derived in Sect.\\ts 4, expressed in terms of a set of galaxy positions. From these expressions, the covariances can be determined for an arbitrary survey geometry. In a forthcoming paper (Kilbinger et al., in preparation), we shall use the results of Sect.\\ts 4 to design an optimized geometry for a planned cosmic shear survey.  For the case of a filled survey geometry, the ensemble average of these covariances can be further analyzed; using a few approximations, we express in Sect.\\ts 5 the covariances for this case in terms of integrals. The corresponding expressions have been evaluated, for a particular cosmological model, and are illustrated in a set of figures. In Sect.\\ts 6 we derive the covariance for the aperture mass dispersion, which can be expressed simply in terms of the covariances of the correlation functions. The variance of the aperture mass dispersion, as well as the covariance, is then explicitly calculated for a survey with filled geometry, showing that indeed the aperture mass at two angular scales decorrelates quickly as the scale ratio decreases away from unity.  We then turn in Sect.\\ts 7 to a simple estimator of the power spectrum of the projected cosmic density field, which can be expresed in terms of the correlation functions. Since the correlation functions will be known only over a finite range in angular separation, the simple estimator we derive is biased. We show that, provided the angular range on which the correlation functions can be measured is as large as can be expected with the next generation of cosmic shear surveys, this bias is indeed {\\it very} small over a large range of wave numbers. We derive the covariance of the power spectrum estimator and calculate it explicitly for the filled survey geometry case; the resulting error bars on the estimated power spectrum are quite a bit smaller than one might have expected, given the simplicity of the approach. In Sect.\\ts 8 we consider the accuracy with which the parameters of the cosmological model can be constrained, given a survey area. In fact, by fitting the correlation function directly to model predictions, even the currently available cosmic shear surveys can yield fairly accurate constraints on cosmological parameter. Finally, we summarize our results in Sect.\\ 9.  In this paper we shall assume that the observable shear is due to the tidal gravitational field of the cosmological matter distribution only; in this case, the two shear components are not mutually independent. This is due to the fact that the gravitational field is a gradient field. In the language of some recent papers (e.g., C02; Pen et al.\\ 2002; SvWM), we thus assume that the shear field is a pure E-mode field. B-modes, (or the `curl component'), can in principle be generated if the intrinsic orientation of the galaxies from which the shear is measured are correlated, e.g. due to tidal interactions of dark matter halos in which these galaxies are formed (Croft \\& Metzler 2000; Pen et al. 2000; Heavens et al. 2001; Catelan et al. 2001; Mackey et al. 2002; Brown et al. 2002). Also, the clustering of source galaxies in redshift space generates a B-mode contribution which, however, turns out to be fairly small (SvWM).  This restriction to E-modes only affects the interrelations between various two-point statistics. Inclusion of B-modes would not change the results of Sects.\\ 3 through 5, and much of Sects.\\ 6 and 7 will also be left unaffected in the presence of a B-mode contribution; we shall indicate this in due course. ",
        "conclusions": "In this paper we have obtained general expressions for the covariance of an estimator for the shear correlation function as it is determined from cosmic shear data. Using the approximation that the four-point function of the shear separates in products of two-point function, the covariance can be expressed directly in terms of the correlation functions, as given in (\\ref{eq:Cpp}--\\ref{eq:Cpm}) and can, for a given data set, be calculated directly. The covariance of the correlation functions depends on the number of pairs that enter their estimate, which in turn depends on the solid angle covered by the survey and the survey geometry (see also Kaiser 1998); in addition, it depends on the intrinsic galaxy ellipticities and the number density of galaxies.  Next, considering a survey geometry of a single compact region of solid angle $A$, we have calculated the ensemble average of the covariances, using approximations which a valid for separations $\\ll\\sqrt{A}$. The ensemble average of the covariances can then be reduced to integrals which are readily evaluated numerically. The estimate for the correlation function $\\xi_-(\\vt)$ decorrelates quickly, i.e. estimates of $\\xi_-$ for two angular scales which differ by more than a factor $\\sim 2$ are essentially decorrelated. On the other hand, the estimates of $\\xi_+$ are correlated over much larger angular scales. The cross-correlation between $\\xi_+(\\vt_1)$ and $\\xi_-(\\vt_2)$ is significant for $\\vt_1\\lesssim \\vt_2$, which is due to the properties of the different filters with which these correlation functions are related to the power spectrum $P_\\kappa(\\ell)$.  Using these ensemble-averaged covariances for the correlation functions, we have obtained the covariances for other two-point measures of the cosmic shear, primarily the aperture mass dispersion and the power spectrum. Of particular interest is the reconstruction of the power spectrum $P_\\kappa(\\ell)$ from the correlation functions; we have constructed a simple estimator for $P_\\kappa$ and the band powers of it in terms of the $\\xi$'s and found that the band power can be obtained with surprisingly large accuracy from even a moderately-sized cosmic shear survey. Finally, we have investigated the confidence regions for the most relevant cosmological parameters ($\\Omega_{\\rm m}$, $\\sigma_8$, $\\Gamma$ and $z_{\\rm s}$) with a maximum likelihood approach. We studied our ability to constrain simultaneously these parameters from a measurement of the shear correlation function, as well as the effect of some level of lack of knowledge using the marginalization technique.  In a future paper, we shall investigate strategies for conducting cosmic shear surveys by optimizing the survey geometry."
    },
    "0206/astro-ph0206461_arXiv.txt": {
        "abstract": "Turbulence may have been produced in the early universe during several kind of non-equilibrium processes. Periods of cosmic turbulence may have left a detectable relic in the form of stochastic backgrounds of gravitational waves. In this paper we derive general expressions for the power spectrum of the expected signal. Extending previous works on the subject, we take into account the effects of a continuous energy injection power and of magnetic fields. Both effects lead to considerable deviations from the Kolmogorov turbulence spectrum. We applied our results to determine the spectrum of gravity waves which may have been produced by neutrino inhomogeneous diffusion and by a first order phase transition. We show that in both cases the expected signal may be in the sensitivity range of LISA. ",
        "introduction": "The high isotropy of the cosmic microwave background radiation (CMBR) testifies that the universe was a quiet child at photon decoupling time. This does not prevent it, however, to have been a quite {\\em turbulent} baby!  Several violent and interesting phenomena may have taken place before the universe became matter dominated. Phase transitions, and re-heating at the end of inflation are  examples (we will discuss another one below) of processes which may have taken the universe through a phase during which thermal equilibrium and homogeneity can be temporarily lost. Turbulence is expected to have developed under such conditions due to the huge value of the Reynolds number. It would be extremely interesting for cosmologists to detect any observable relic of such turbulent periods as this may shed light on the early stages of the universe and on fundamental physics which is not yet probed in the laboratories. Unfortunately, due to the finite thickness of the last scattering surface, density or metric perturbations which may have been produced by any causal process  before the matter radiation equality time than can hardly give rise to detectable imprints on the CMBR. The detection of cosmological background(s) of gravitational waves (GWs) may offer the only possibility to probe turbulence in the remote past.  Although any direct detection of GWs is still missing,  this subject is not any more purely academic as a new class of ground and space based observatories of GWs are under construction or in advanced project. The amplitude and frequency sensitivity of some of these instruments are in the proper range to probe many  interesting astrophysical and cosmological processes \\cite{maggiore}. The most interesting project from our point of view is LISA (Laser Interferometer Space Antenna) which is scheduled to be launched around the end of this decade and will possibly achieve a sensitivity in $h_0^2 \\Omega_{\\rm gw}$ of $10^{-12}$ at milliHertz frequencies \\cite{LISA}.  So far the interest on GWs produced by cosmological turbulence has been mainly stimulated by the possibility that turbulence could have been generated at the end of the electroweak phase transition (EWPT) \\cite{KKT,KMK,alberto}. The idea is that, if the EWPT is first order, turbulence should have been produced when the expanding walls of bubbles containing the broken symmetry phase, or the shock waves preceding them, collide at the end of the transition. As discussed in \\cite{KMK,alberto} the stochastic background of GWs produced during this process may be detectable by LISA.  Recently two of us \\cite{dario} noticed that another, possibly detectable, background of GWs  may have been produced before neutrino decoupling, that is much later  than the EWPT. This mechanism requires that the net lepton number density ($N_a({\\bf x}) \\equiv n_{\\nu_a}({\\bf x}) - n_{\\bar \\nu_a}({\\bf x})$,\\ \\ $a = e,\\mu,\\tau$) of one, or more, neutrino species was not uniform and changed in space over some characterist scale which was smaller than the Hubble horizon at the neutrino decoupling time. A net flux of neutrinos should then be produced along the lepton number gradients when the neutrino mean free path became comparable to the size of the fluctuations in the neutrino number. It was shown that, depending on the amplitude of the fluctuations, such a flux of neutrinos may be able to stir chaotically the cosmic plasma producing magnetic fields and GWs.  Fluctuations in the neutrino number, as those which are required to power all this rich set of effects, can be a by-product of leptogenesis at the GUT scale (this is possible for example in the Affleck-Dine \\cite{DolKir} scenario  of leptogenesis) or be a consequence of active-sterile resonant neutrino conversion before neutrino decoupling \\cite{DiBari}.  The aim of this paper is to estimate the characterists amplitude and spectral distribution of GWs produced  by the mechanism discussed in \\cite{dario} and during a first order phase transition. Our treatment follows that of Kosowsky, Mack and Kahniashvili \\cite{KMK} (KMK) and extend it by introducing some physically relevant generalizations.  In Sec.~\\ref{turbo} we will show how the Kolmogorov spectrum of turbulence (which is that adopted by KMK) is modified when energy injection into the primordial plasma takes place over a continuos range of scales. This is motivated by the observation that  neutrino number fluctuations, which might power turbulence, generally will not have a unique size but have a more complex spectral distribution. The same consideration can be naturally applied to other kind of stochastic sources. Since strong magnetic fields are expected to be produced by the mechanism presented in Ref.~\\cite{dario} and during first order phase transitions \\cite{Baym,report}, in Sec.~\\ref{MHD} we will also investigate possible Magneto-Hydro-Dynamical (MHD) effects on the turbulent spectrum. We will see that these kind of effects may give rise to quite drastic consequences on the expected GW signal. In Sec.~\\ref{GWs} we apply the results of the previous two sections to determine the general expression of the GW power spectrum produced by cosmological turbulence. In Sec.~\\ref{neutrino} we estimate the signal to be expected at LISA  because of GWs generated by neutrino inhomogenoeus diffusion. In Sec.~\\ref{EWPT} we will determine the characteristics of the GW signal produced by a first order phase transition paying particular attention to the case of the EWPT. We will show that MHD effects may play a crucial role by enhancing the expected signal. Finally Sec.~\\ref{end} contains our conclusions. ",
        "conclusions": "\\label{end}  In this paper we investigated the characteristics of the GW signal sourced by turbulence in the early universe. We extended previous work on the subject by considering the effects on the turbulence energy spectrum of a continuous stirring power and of magnetic fields.  In the  first part of our work we derived a suitable parameterization for the turbulent energy spectrum in the case of a continuous stirring power with and without external magnetic fields. Interestingly we found that the MHD turbulent spectrum can be written in the same form as the Navier-Stokes ones upon a simple rescaling of only two parameters. Our treatment is original and may find applications which go beyond the aims of this paper. Following the approach of Kosowsky {\\em et al.}~\\cite{KMK}, who consider only Kolmogorov type of turbulence, we used our formulation to determine the characteristics of the GW signal produced by cosmological turbulence with a more general energy spectrum. In the Kolmogorov case we found minor discrepancies with the result reported in \\cite{KMK}. Deviations from the Kolmogorov spectrum may have relevant consequences for the GW signal. Indeed, since the modified spectra are generally less steep than the Kolmogorov's, these deviations will result in a stronger signal at high frequencies. Furthermore, our analysis may allow to extract valuable informations about the nature of the turbulence source and the presence of primordial magnetic fields from the GW background power spectrum which may be measured by forthcoming experiments.  We applied our results to estimate the GW expected signal for two possible kind of sources. In sec.~\\ref{neutrino} we considered GW production by neutrino inhomogeneous diffusion according to the mechanism proposed in Ref.~\\cite{dario}. We showed that a detectable signal can be produced only if the amplitude of the lepton number fluctuations is close to unity over a wide range of wavenumbers. This possibility is not unreasonable as active-sterile neutrino oscillations or some models of leptogenesis based on the Affleck-Dine scenario can indeed give rise to domains with opposite lepton number. Although we argued that the former scenarios can hardly give rise to GW signal detectable by LISA, the latter scenario offers more promising observational perspectives.  Finally in sec.~\\ref{EWPT} we considered GW production by turbulence at the end of a first order phase transition. In the absence of strong magnetic fields no substantial deviations have to be expected from the results of Ref.~\\cite{KMK} based on the assumption of a Kolmogorov turbulence spectrum. In the case of the EWPT  the GW signal may be above the LISA planned sensitivity only if the transitions is very strong with large bubble wall velocities. This may be achieved in next to minimal extensions of the supersymmetric standard model \\cite{alberto}. A more favorable situations may be obtained if strong magnetic fields were present at the end of the transition. This is actually to be expected in the case of the EWPT if it is first order \\cite{Baym}. We showed that in this case the signal can be strongly enhanced with respect to the non-magnetic case and be detectable by LISA if bubble wall velocity was not too much smaller than unity. Similar consideration may apply to other physical situations like, for example, during reheating at the end of inflation. These results open new perspectives for a successful detection of GWs backgrounds produced in the early universe.  \\subsection*"
    },
    "0206/hep-ph0206113_arXiv.txt": {
        "abstract": "The dependence of the nuclear force on standard model parameters plays an important role in bounding time and space variations of fundamental couplings over cosmological time scales. We discuss the quark-mass dependence of deuteron and di-neutron binding in a systematic chiral expansion. The leading quark-mass dependence of the nuclear force arises from one-pion exchange and from local quark-mass dependent four-nucleon operators with coefficients that are presently unknown. By varying these coefficients while leaving nuclear observables at the physical values of the quark masses invariant, we find scenarios where two-nucleon physics depends both weakly and strongly on the quark masses.  While the determination of these coefficients is an exciting future opportunity for lattice QCD, we conclude that, at present, bounds on time and space variations of fundamental parameters from the two-nucleon sector are much weaker than previously claimed. This brings into question the reliability of coupling-constant bounds derived from more complex nuclei and nuclear processes. ",
        "introduction": "The recent observation suggesting that the fine-structure constant was smaller in the past~\\cite{dalpha} than it is today has led to renewed interest in the idea of using time (and space) variation of fundamental parameters as a probe of high-energy physics.  Based on the principle that all that is not forbidden is mandatory, time variation in the electromagnetic force is perhaps not surprising. Moreover, one would expect that the other forces of nature vary with time as well.  This idea is quantifiable by assuming a suitable grand unified theory (GUT) at a scale $M_{\\sss GUT}$ or a particular brane-world scenario~\\cite{CF,matt,Chacko:2002mf}. The assumption that the short-distance couplings of the theory are related at all times leads to relations between the variation of the strength of the electromagnetic interaction, $\\alpha_{\\rm em}$, and the variation of other standard model parameters, such as the light quark masses, $m_q$.  As a change in $m_q$ and $\\alpha_{\\rm em}$ will naively lead to a change in the positions of nuclear energy levels, special interest has been paid to the light-element abundances predicted by big bang nucleosynthesis (BBN) and also to the abundance of isotopes produced by the Oklo ``natural reactor'' in the hope that these abundances can be used to constrain high-energy physics~\\cite{CF,matt,KPW,Barrow,Dixit,campbell,Damour,agrawal,hummer,Fujii,Shera,BIR,Csoto,Chiba,dent,fair,Dolgov,IK,shuryak,olive,uzan}.  In this work we will critically analyze the two-nucleon sector using an effective field theory (EFT) that respects the approximate $SU(2)_L\\otimes SU(2)_R$ chiral symmetry of QCD and has consistent power-counting~\\cite{We90,KSWa,KSWb,Be01}.  This theory allows for a systematic study, consistent with QCD, of the bound-state in the $\\siii-\\diii$ coupled-channels, the deuteron, and the scattering amplitude in the $\\si$ channel. This is the first study of the $m_q$-dependence of these quantities that is {\\it complete} at next-to-leading order (NLO) in a consistent EFT. Unfortunately, due to the current lack of understanding of the low-energy behavior of QCD, only weak bounds, if any, can be placed on the  time-dependence of fundamental couplings from the two-nucleon sector. We will explicitly demonstrate this by finding plausible values for couplings in the EFT for which the deuteron binding energy varies relatively little, and for which the di-neutron system in the $\\si$ channel is unbound, over a relatively wide range of quark masses.  The deuteron is an interesting object to study for a variety of reasons. From a fundamental point of view it is quite intriguing as it is bound by only $B_{\\rm d} = 2.224644\\pm 0.000034~{\\rm  MeV}$, which is much smaller than the typical scale of strong interactions, $\\Lambda_{\\sss QCD}$. Clearly, in order to arrive at such a small binding energy, fine-tunings are involved~\\footnote{In the pionless EFT one finds that both the $\\si$ and the $\\siii -\\diii$ coupled-channels are very close to an unstable IR fixed point of the renormalization group. At this fixed point, the scattering lengths are infinite~\\cite{KSWb,Birse}.} and consequently naive dimensional analysis (NDA) as applied to the single-nucleon sector or the mesonic sector, which we will call naive naive dimensional analysis (N$^2$DA), on this system is doomed to fail.  Unfortunately, all previous attempts to extract bounds on the variations of fundamental couplings from the deuteron, and in general the two-nucleon sector, have implemented N$^2$DA. From a more phenomenological standpoint, the smallness of $B_{\\rm d}$ plays a key role in the synthesis of light elements in BBN. The impressive agreement between the predictions of BBN and observation suggests that new physics that would have significantly modified the deuteron at the time of BBN is absent. The $\\si$ channel is quite similar to the $\\siii-\\diii$ coupled-channels in one important way, its scattering length is unnaturally large, $a^{(\\si)}=-23.714\\pm 0.013~{\\rm fm}$. While there is no bound state in this channel for the physical values of the light-quark masses, a small increase in the strength of the nuclear force would bind two nucleons in this channel. The existence of a bound state in this channel, e.g. a di-neutron, $nn$, in the nucleosynthesis epoch would be quite profound and could substantially modify the predictions of BBN. ",
        "conclusions": "We have explored the light-quark mass dependence of low-energy nucleon-nucleon interactions.  The motivation for this work was to determine if, in fact, bounds could be set on the time-variation of fundamental couplings from nuclear processes, such as those occurring during big bang nucleosynthesis.  We have demonstrated the existence of sets of strong interaction couplings in the low-energy effective field theory describing the nucleon-nucleon interaction that are consistent with all available data and with naive dimensional analysis for which the di-neutron remains unbound, and the deuteron remains loosely bound over a wide range of light quark masses.  We do not mean to imply that these are the sets of couplings that nature has chosen, but rather that this scenario is not excluded at present.  Thus, we conclude that bounds that have previously been set on the time-variation of fundamental couplings from processes in the two-nucleon sector are not rigorous and should be discarded.  Our calculation does suffer from some limitations.  First, we have not included electromagnetism, and thus could not address the $pp$-system in the $\\si$-channel.  However, we do not believe that its inclusion will modify the results we have presented in any significant way. The same we believe to be true of isospin breaking. Second, and perhaps the most important limitation, is that we have not included the strange quark in our discussions, and have worked with $SU(2)_L\\otimes SU(2)_R$ chiral symmetry. Variations in the strange (and charm, bottom and top) quark mass will manifest themselves as changes in the values of the coefficients in the Lagrange density, $C_i$ and $D_i$.  At this point in time the tools do not exist to analyze this scenario.  For instance, one might consider developing an EFT with three light-quark flavors. However, since nuclei are such finely-tuned systems, and $SU(3)$-breaking is not so small, it is likely that a perturbative treatment of nuclei using the approximate $SU(3)_L\\otimes SU(3)_R$ chiral symmetry of QCD would converge very slowly, if at all.  From a conventional nuclear physics point of view, the $m_q$-dependence of the nuclear potentials not only requires knowledge of the $m_q$-dependence of one-pion exchange, but also of the $m_q$-dependence of the $\\rho$, $\\omega$, $\\phi$,...  masses and their couplings to nucleons. This would appear to be an intractable problem. It is possible that some of the traditional treatments of the nucleon-nucleon interaction~\\cite{Nijmegen,WSS,Pudliner:1997ck,Timmermans}, and treatments of light nuclei could mimic these dependences by variations in their ad-hoc short-distance components of the nucleon-nucleon and higher-body potentials.  The recently developed effective field theory tools organize this problem in a very simple way, and thus at NLO in the EFT expansion, only the $D_i$ are needed. In experimental processes examined to date that, in principle, allow one to separate the $D_i$ from the $C_i$, it is found that there are other amplitudes contributing to the process that dominate over the $D_i$ contributions, rendering such a separation exceedingly difficult. In the absence of an experimental determination of the $D_i$, it would appear that the only viable means for determination of this $m_q$-dependence is through lattice QCD simulations~\\footnote{A first attempt at computing nucleon-nucleon scattering lengths in quenched lattice QCD has been made in Ref.~\\cite{qlattice}. A discussion of the two-nucleon potential in quenched and partially-quenched QCD can be found in Ref.~\\cite{BSpot}. For a recent discussion, see Ref.~\\cite{richards}.}. We find this to be very strong motivation to pursue a lattice QCD program focused on the two-nucleon sector.  \\vskip0.2in \\noindent {\\large\\bf Acknowledgements}  \\noindent We thank George Fuller for useful discussions and Iain Stewart for helpful comments on the manuscript. We are especially grateful to Evgeni Epelbaum for pointing out an error in an earlier version of the manuscript, and for useful discussions. This research was supported in part by the DOE grant DE-FG03-97ER41014.  \\vspace{1cm}"
    },
    "0206/astro-ph0206083_arXiv.txt": {
        "abstract": "{The location of the red edge of the ZZ Ceti instability strip is defined observationally as being the lowest temperature for which a white dwarf with a H-rich atmosphere (DA) is known to exhibit periodic brightness variations. Whether this cut-off in flux variations is actually due to a cessation of pulsation or merely due to the attenuation of any variations by the convection zone, rendering them invisible, is not clear. The latter is a theoretical possibility because with decreasing effective temperature, the emergent flux variations become an ever smaller fraction of the amplitude of the flux variations in the interior. In contrast to the flux variations, the visibility of the velocity variations associated with the pulsations is not thought to be similarly affected. Thus, models imply that were it still pulsating, a white dwarf just below the observed red edge should show velocity variations. In order to test this possibility, we used time-resolved spectra of three DA white dwarfs that do not show photometric variability, but which have derived temperatures only slightly lower than the coolest ZZ Ceti variables. We find that none of our three targets show significant periodic velocity variations, and set 95\\% confidence limits on amplitudes of 3.0, 5.2, and 8.8\\,km\\,s$^{-1}$. Thus, for two out of our three objects, we can rule out velocity variations as large as 5.4\\,\\kms\\ observed for the strongest mode in the cool white dwarf pulsator \\object{ZZ Psc}. In order to verify our procedures, we also examined similar data of a known ZZ Ceti, \\object{HL Tau 76}. Applying external information from the light curve, we detect significant velocity variations for this object with amplitudes of up to 4\\,km\\,s$^{-1}$. Our results suggest that substantial numbers of pulsators having large velocity amplitudes do not exist below the observed photometric red edge and that the latter probably reflects a real termination of pulsations. ",
        "introduction": "\\label{sec:intro} In spite of advances in theoretical formulations and observational capabilities, the details of the inner workings of the cool H-rich pulsating white dwarfs (DAVs or ZZ Cetis) continue to elude us.  Early attempts to explain the nature of the driving of the pulsations of ZZ Cetis postulated that driving occurred via the $\\kappa$-mechanism -- i.e. in a manner akin to that in Cepheids and $\\delta$-Scuti stars -- \\citep[e.g.][]{danv:81,dzak:81,wing:82} thereby ignoring the effects of pulsation on the convection zone (also known as the ``frozen-in'' approximation) even though radiative flux transport is negligible in these regions. \\citet{brick:91} realised that the response of the convection zone to the perturbation is almost instantaneous ($\\sim 1$\\,s) compared to the periods of the g-modes (hundreds of seconds). He found that as a result, the convection zone itself can drive the pulsations: part of the flux perturbations entering the convection zone from the largely radiative interior are absorbed by the convection zone and are released half a cycle later. This interaction between the perturbation and the response of the convection zone drives the g-modes.  Within the framework of this ``convective-driving'' mechanism, recently confirmed analytically by \\citet{gw:99a}, the blue edge of the instability strip i.e. the hottest temperature for which pulsations are excited and at which pulsations are expected to be observable, is set by the condition that $\\omega \\tau_{c} \\approx 1$ for radial order, {\\em n \\/} = 1 and spherical degree $\\ell=1$; $\\omega$ is the radian frequency and $\\tau_{c}$ the thermal time constant of the convection zone. \\citep[Note that $\\tau_{c}$ is different from the `global' thermal time scale, $t_{t}$ in \\citet{brick:91}. Where necessary, we follow the notation of][] {gw:99a}.  The location of the red edge is less clearly defined. As long as a mode is driven, its intrinsic amplitude is likely to remain roughly constant as it is most likely set by parametric resonance with stable daughter modes \\citep{wg:01}. As the white dwarf cools, the depth of the convection zone increases, and damping in the shear layer at the base of the convection zone becomes stronger. At some stage, damping will exceed the driving due to the convection zone. This sets the physical red edge, beyond which pulsations are no longer excited.  Observationally, though, the red edge may appear at {\\em higher\\/} temperatures. This is because, as the convection zone deepens, the amplitude of flux variations at the photosphere becomes an ever smaller fraction of the amplitude in deeper layers; the convection zone acts as a low-pass, frequency-dependent filter leading to a reduction in flux given by \\citep{gw:99a} \\begin{equation} \\left(\\frac{\\Delta F}{F}\\right)_{\\rm ph} = \\left(\\frac{\\Delta F}{F}\\right)_{\\rm b}\\frac{1} {\\sqrt{1 + (\\omega\\tau_{\\rm c})^{2}}} \\label{eq:dfof} \\end{equation} where the subscripts ``ph'' and ``b'' refer to the photospheric flux variation and that at the base of the convection zone respectively, for a particular eigenmode. From Eq. (\\ref{eq:dfof}) one expects that the observed amplitude of a given mode will decrease smoothly with increasing $\\tau_{\\rm c}$, and thus with decreasing temperature. This might seem inconsistent with the observations, which show a rather sharply defined red edge. However, $\\tau_{\\rm c}$ depends very steeply on temperature, and hence it may also be that the observations do not yet resolve the transition from easily to barely visible \\citep[see Fig. 8 in][]{wg:99}.  Given the uncertainties in the physical processes leading to damping, it is not clear whether the observed red edge corresponds to the physical red edge, or whether it is just an apparent red edge where pulsations are still driven, but do not give rise to observable flux variations at the photosphere. The implication, then, is that a perfectly constant white dwarf might actually still be pulsating!  Theoretical uncertainties aside, there are also uncertainties in interpreting the observations. Indeed, over the years, the location and extent of the observed ZZ Ceti instability strip have undergone several transformations \\citep[e.g.,][]{greens:82,robetc:95,giov:98}. This is due in part to the difficulty in accurately determining the atmospheric parameters of objects having convectively unstable atmospheres, as some version of the mixing length prescription usually has to be assumed. Additionally, at optical wavelengths, the Balmer lines attain their maximum strengths in the middle of the instability strip. The unfortunate consequence is that varying the atmospheric parameters gives rise to only slight changes in the appearance of the spectra, making it difficult to uniquely fit observed spectra \\citep[e.g.,][]{kv:96}. Supplementary constraints in the form of uv spectra, parallaxes, gravitational redshifts, have almost become a prerequisite.  Most studies to date have only considered flux variations. However, velocity variations are necessarily associated with these flux variations and though small -- of the order of a few \\kms\\ -- have nevertheless been measured in ZZ Ceti white dwarfs: securely in ZZ Psc \\citep{vkcw:00}, and probably also in \\object{HS 0507+0434B} \\citep{kotak:02}. That there are negligible vertical velocity gradients in the convection zone due to damping by turbulent viscosity is a central tenet of the convective-driving theory \\citep{brick:90,gw:99b}. Indeed for shallow convection zones, this simplification has been shown to hold in the 2D hydrodynamical simulations of \\citet{gautschy:96}. This means that the horizontal velocity is nearly {\\em independent} of depth in the convection zone so that although photometric variations become difficult to detect around the red edge of the instability strip as mentioned above, velocity variations pass virtually undiminished through the convection zone.  \\begin{table*}[htb] \\caption[]{Observations} \\fontsize{8.5}{10}\\selectfont \\label{tab:obs} \\begin{centering} \\begin{tabular}{lcccccccccc} \\hline & WD  &  \\Teff  & $\\log g$ & V & Start & End & Exposure & No. of & Scatter & Measurement \\\\ & number &  (K)    &          &   & U.T.  & U.T. & Time (s) & frames & (\\kms)  & error (\\kms)  \\\\ \\hline HL Tau 76 (s)& 0416+272 &   11440  & 7.89 $\\pm$ 0.05 & 15.2 & 12:58:55 & 14:59:06 & 20  &   213    & 9.0  & 6.7  \\\\ G 1-7 (s)    & 0033+013 &   11214  & 8.70 $\\pm$ 0.06 & 15.5 & 08:37:49 & 10:39:50 & 74  &  \\phn83  & 9.8  & 7.7 \\\\ G 67-23 (s)  & 2246+223 &   10770  & 8.78 $\\pm$ 0.03 & 14.4 & 07:33:57 & 08:48:15 & 74  &  \\phn50  & 3.5  & 2.4  \\\\ G 126-18     & 2136+229 &   10550  & 8.17 $\\pm$ 0.05 & 15.3 & 04:46:06 & 06:27:25 & 30  &   136    & 1.4  & 7.6 \\\\ G 126-18 (s) &          &          &                 &      & 06:58:18 & 08:58:03 & 74  &  \\phn80  & 5.2  & 4.0  \\\\ \\hline \\end{tabular} \\tablecomments{The effective temperature and $\\log g$ for each of the 3 red edge objects are taken from \\citet{kep:95} and are inferred from optical data only using the ML2, $\\alpha$=1 prescription for describing the convective flux. For HL Tau 76, these quantities are taken from \\citet{berg:95} and are constrained using both uv and optical spectra for ML2/$\\alpha$=0.6. The scatter due to wander is calculated at a representative wavelength of 4341\\,{\\AA} i.e. at H$\\gamma$. The measurement error is the typical internal error on the line-of-sight velocity associated with the fits to the line profiles as described in Sect. \\ref{sec:vels}. ``(s)'' refers to the Oct. 1997 service run data. The number of frames refers to the number of {\\em useful\\/} frames.} \\end{centering} \\end{table*}  Observationally, the consequence would be that the putative, perfectly constant white dwarf might reveal its pulsating nature by velocity variations. If objects below the red edge follow the same trend as the known pulsators i.e. longer pulsation periods and higher amplitudes with decreasing effective temperature \\citep{clem:93}, then we expect the highest velocity amplitudes at the longer ($\\ga$ 600\\,s) periods. For ZZ Psc (a.k.a. G 29-38), a well-known pulsator close to the red edge, \\cite{vkcw:00} measured a velocity amplitude of 5.4\\,\\kms\\ for the strongest mode at 614\\,s. For all our objects, which are cooler, we expect velocity amplitudes at least as large as those measured for ZZ Psc -- lower values would constitute a non-detection.  Armed with this testable theoretical expectation, we look for variations in the line-of-sight velocities in objects that lie just below the photometric red edge of the instability strip. Our three targets are chosen -- subject to visibility constraints during the scheduled run -- from the list of \\citet{kep:95} who find a handful of non-pulsating ZZ Cetis close to the red edge of the instability strip. \\citet{kep:95} inferred relatively high masses for two of our three targets (see Table \\ref{tab:obs}). This inference, though dependent on the assumed convective prescription, would imply that these objects have instability strips at higher temperatures.  In addition to the above, we search for line-of-sight velocity variations in the first ZZ Ceti to be discovered, HL Tau 76 \\citep{land:68}, using exactly the same instrumental setup and reduction techniques so that it serves as a comparison case. Treating it subsequently as a ZZ Ceti variable (Appendix \\ref{sec:hlt}) allows us to indirectly constrain the spherical degree of the eigenmodes. ",
        "conclusions": "We have analysed velocity curves of three white dwarfs below the photometric red edge of the instability strip to check whether the observed red edge corresponds to the physical one. For ZZ Psc, a white dwarf close to the red edge but still within the instability strip, \\citet{vkcw:00} found a maximum velocity amplitude of $5.4 \\pm 0.8\\,$\\kms. We had expected larger modulations but for two objects we can exclude the presence of signals as strong as those in ZZ Psc with high confidence; for one of these two, G 67-23, we can exclude the presence of any signal greater than 3\\,\\kms.  The above indicates that pulsations actually cease below the observed red edge. How secure is this conclusion? We see three weaknesses. The first is that in principle, destructive beating between different components in a multiplet might conceal a real signal in our short data sets. This, however, would be unlikely if multiple real signals were present. Furthermore, two data sets for the same object (G\\,126-18) taken during different observing runs are unlikely to be adversely affected. Two of the modulations found in G 126-18 (viz. 220\\,s and 269\\,s) seemed to be consistent with being at the same frequency and of the same amplitude -- within the errors -- for the two different observing runs. We carried out further Monte Carlo simulations to determine the likelihood of finding relatively strong peaks at the same frequency in two separate data sets. On the basis of these simulations, we conclude that these peak coincidences are almost certainly due to chance: even for the best case, the peaks at 265\\,s and 269\\,s, there is a 7\\% probability of a chance coincidence.  The second weakness is that our sensitivity turned out to be more marginal than we had hoped. We included one known ZZ Ceti, HL Tau 76, in the sample.  Treating it as a red-edge object afforded a check on our methods. However, even for this bona-fide pulsator, we found that the velocity curve on its own was not sufficient to demonstrate velocity variations. A hint of a real signal was present, but this could only be confirmed by imposing external information from the light curve, information we lacked for the red-edge objects. In the process, HL Tau 76 was added to the select group of ZZ Cetis for which velocity variations have been detected.  The above might suggest that our experiment was simply not sensitive enough. One has to keep in mind however, that for the different objects rather different sensitivities were reached. For one object, G 67-23, we can exclude the presence of velocity amplitudes as small as 3\\,\\kms\\ i.e. even smaller than the 4\\,\\kms\\ peak corresponding to strongest mode in HL Tau 76. Furthermore, on theoretical grounds alone, a velocity amplitude somewhat smaller than that of ZZ Psc is expected for HL Tau 76 given that it has a slightly higher temperature and its strongest mode, a slightly shorter period. This expectation is supported by the fact that the velocity amplitude of the dominant mode is in turn larger than that found for HS~0507+0434B, which is somewhat hotter still and has an even shorter dominant periodicity. On the same grounds, one would expect, as mentioned earlier, the red-edge objects to have larger velocity amplitudes than those observed in ZZ~Psc, which, if present, would have been seen in two of the three objects.\\footnote{We remind the reader that the amplitude which a driven mode reaches is not expected to depend on temperature, but on resonance conditions with daughter modes; see \\citet{wg:01}.}  The third weakness is perhaps the most severe: G\\,1-7 and G\\,67-23 have relatively high inferred masses, which implies that these objects would have ceased to be pulsationally active at a higher effective temperature than that expected for a typical $0.6\\,M_\\odot$ white dwarf, and are therefore not that close anymore to the red edge. At the same time, G\\,126-18, which has a normal mass, is the coolest of our three observed targets. Note though that the derived effective temperatures and surface gravities depend on the treatment of convection.  In summary, while not ironclad, our results indicate that pulsations have actually ceased below the observed red edge rather than having become photometrically undectable, and that the theoretical expectation of an extended instability strip, beyond the observed red edge is flawed. In order to settle the issue observationally, more sensitive measurements on objects closer to the red edge would be worthwhile; theoretically, detailed hydrodynamic modelling might go some way towards understanding the interplay of the various physical processes in defining the red edge."
    },
    "0206/astro-ph0206426_arXiv.txt": {
        "abstract": "We present observations of the dwarf novae GW Lib, V844 Her, and DI UMa.  Radial velocities of H$\\alpha$ yield orbital periods of $0.05332  \\pm 0.00002$ d (= 76.78 m) for GW Lib and and $0.054643 \\pm 0.000007$ d (= 78.69 m) for V844 Her. Recently, the orbital period of DI UMa was found to be only $0.054564 \\pm 0.000002$ d (= 78.57 m) by \\citet{fr99}, so these are the three shortest orbital periods among dwarf novae with normal-abundance secondaries.  GW Lib has attracted attention as a cataclysmic binary showing apparent ZZ Ceti-type pulsations of the white dwarf primary.  Its spectrum shows sharp Balmer emission flanked by strong, broad Balmer absorption, indicating a dominant contribution by white-dwarf light. Analysis of the Balmer absorption profiles is complicated by the unknown residual accretion luminosity and lack of coverage of the high Balmer lines.  Our best-fit model atmospheres are marginally hotter than the ZZ Ceti instability strip, in rough agreement with recent ultraviolet results from HST.  The spectrum and outburst behavior of GW Lib make it a near twin of WZ Sge, and we estimate it to have a quiescent $M_V \\sim 12$.  Comparison with archival data reveals proper motion of $65 \\pm 12$ mas yr$^{-1}$.  The mean spectrum of V844 Her is typical of SU UMa dwarf novae.  We detected superhumps in the 1997 May superoutburst with $P_{\\rm sh} = 0.05597 \\pm 0.00005$ d. The spectrum of DI UMa appears normal for a dwarf nova near minimum light.  These three dwarf novae have nearly identical short periods but completely dissimilar outburst characteristics. We discuss possible implications. ",
        "introduction": "Cataclysmic variable stars are close binary systems in which mass is transferred from a low-mass secondary onto a white dwarf; \\citet{warn} has written an excellent monograph on cataclysmics. In the dwarf nova subclass, the transferred matter is thought to accumulate in an accretion disk until a critical density is reached, whereupon the transfer of matter through the disk dramatically increases, causing a brightening of the star.  The outbursts of dwarf novae range widely in amplitude and frequency, with the less-frequently outbursting systems tending to have the greatest amplitudes (the Kukarkin-Parenago relation).  The orbital periods $P_{\\rm orb}$ of dwarf novae range from $\\sim 12$ hours down to about 75 minutes, near the theoretical minimum period for a hydrogen-rich donor star \\citep{rjw82}. (We do not here consider systems with helium donor stars, which can reach much shorter periods.)  Dwarf novae with $P_{\\rm orb}$ less than about 3 h (effectively, those shortward of the 2-3 h period `gap') show occasional large-amplitude, long-duration outbursts, called superoutbursts, which are accompanied by photometric oscillations, called superhumps, which have periods $P_{\\rm sh}$ a few percent {\\it longer} than $P_{\\rm orb}$.  Superoutbursters are classified as SU UMa stars, after their prototype. But SU UMa stars show a wide range of intervals between outbursts, from about 4 days for ER UMa and its brethren \\citep{jc,p95,rht95,ms95,twoers}, to several decades for WZ Sge and similar stars \\citep{wxcet,alcom,egcnc,katoeruma95}.  Here we report extensive observations of the dwarf novae GW Lib and V844 Her, which prove to have very short orbital periods. We also present a spectrum of DI UMa, another very short period system.  In spite of their similar orbital periods, these three systems have quite different outburst characteristics.  Section 2 of this paper describes observational protocols and the analysis which produced orbital periods for GW Lib and V844 Her.  Sections 3, 4, and 5 contain more detailed discussions of GW Lib, V844 Her, and DI UMa respectively. Finally, Section 6 contrasts the three systems and discusses implications for cataclsymic binary evolution. ",
        "conclusions": "In Table 5 we list the seven dwarf novae (or dwarf-nova candidates) with the shortest well-established orbital periods. Two of the systems have periods shorter than GW Lib.  V485 Cen has a dramatically short 59-min period.  \\citet{aug96} explain this by invoking a low hydrogen abundance for the secondary, so it may not belong among the more usual hydrogen accretors. Support for this idea comes from the apparently related system RX2329+06, which has a K-type secondary at $P_{\\rm orb} = 64$ min, a period at which a CV secondary with normal hydrogen abundance would be be much cooler.  \\citet{thor02} show that a helium-enriched secondary matches the observations, and propose that mass transfer began near the end of core hydrogen burning. Another system, RX2353$-$38, resembles a dwarf nova spectroscopically, and it probably is one, but it has never been observed to erupt.  Thus the stars studied here have the three shortest periods among well-established, hydrogen-accreting dwarf novae.  In the most straightforward scenarios for CV evolution (see \\cite{king88} for a review),  the properties of dwarf novae are regulated by the accretion rate, which is dependent on the mass of the secondary, which depends on orbital period. Thus the outburst properties should depend in a simple way on $P_{\\rm orb}$.  In particular, stars of the shortest orbital period should have very rare outbursts, like the well-studied prototype WZ Sge.  The truth appears to be much more complex.  One of these stars, GW Lib, is indeed an excellent match to WZ Sge: intrinsically very faint, with a recurrence period probably exceeding 10 years.  Yet DI UMa is a rather bright star (at the distance estimated by \\cite{fr99} it would have $M_V = +8.4$ at minimum), and one of the most frantic dwarf novae in the sky, erupting at least 100 times more frequently.  And V844 Her is an intermediate case, but unusual in never having shown any normal outbursts (despite the modest interval of $\\sim 280$ d between superoutbursts).  This is more or less the full range of activity (and brightness) displayed by dwarf novae.  Yet it occurs in three binaries with orbital periods equal to within 2 minutes.  Why this should be so remains a mystery.  Perhaps there are cycles in the mass-transfer rate over timescales long compared to the historical record, but short compared to the evolutionary timescale \\citep{kingmtcycle,wumtcycle}. Also, in the most common picture of evolution at short period, the angular-momentum loss is thought to be dominated by gravitational radiation, leading to a convergence of evolution among different systems and a universal value for the minimum period.  But if another, more idiosyncratic angular momentum loss mechanism persists --- magnetic braking is an obvious candidate --- the minimum period may not be as clearly defined, and the observed diversity of stars near the minimum may be easier to understand. Indeed, \\citet{kingmtcycle} find that mass-transfer cycles among short-period systems are damped out unless there is some small consequential angular momentum loss (that is, an angular momentum loss mechanism for which $\\dot J$ increases modestly with $\\dot M$).  Another line of evidence also suggests an extra channel for angular momentum loss. \\citet{patlate98} notes the apparent discrepancy between the large number of short-period and `post-bounce' systems predicted by theory, and the much smaller number of such systems which have actually been observed; a second angular momentum loss mechanism which destroys short-period systems would help ameliorate that discrepancy, as well.  {\\it Acknowledgments.} We thank the NSF for support through grant AST 9987334, and the MDM staff for their support.  SV is a QEII fellow of the Australian Research Council."
    },
    "0206/astro-ph0206495_arXiv.txt": {
        "abstract": "We have used a combination of stellar population synthesis and photoionization models to develop a set of ionization parameter and abundance diagnostics based only on the use of the strong optical emission lines.  These models are applicable to both extragalactic \\HII\\ regions and star-forming galaxies. We show that, because our techniques solve explicitly for both the ionization parameter and the chemical abundance, the diagnostics presented here are an improvement on earlier techniques based on strong emission-line ratios.  Our techniques are applicable at all metallicities.  In particular, for metallicities above half solar, the ratio \\NIIOII\\ provides a very reliable diagnostic since it is ionization parameter independant and does not have a local maximum.  This ratio has not been used historically because of worries about reddening corrections. However, we show that the use of classical reddening curves is quite sufficient to allow this \\NIIOII\\ diagnostic to be used with confidence as a reliable abundance indicator.  As we had previously shown, the commonly-used  abundance diagnostic \\R23\\ depends strongly on the ionization parameter, while the commonly-used ionization parameter diagnostic \\OIIIOII\\ depends strongly on abundance. The iterative method of solution presented here allows both of these parameters to be obtained without recourse to the use of temperature-sensitive line ratios involving faint emission lines.  We compare three commonly-used abundance diagnostic techniques and show that individually, they contain systematic and random errors. This is a problem affecting many abundance diagnostics and the errors generally have not been properly studied or understood due to the lack of a reliable comparison abundance, except for very low metallicities, where the \\OIII~$\\lambda$4363 auroral line is used.  Here, we show that the average of these techniques provides a fairly reliable comparison abundance indicator against which to test new diagnostic methods.  The cause of the systematic effects are discussed, and we present a new `optimal' abundance diagnostic method based on the use of line ratios involving \\NII, \\OII, \\OIII, \\SII\\ and the Balmer lines. This combined diagnostic appears to suffer no apparent systematic errors, can be used over the entire abundance range and significantly reduces the random error inherent in previous techniques.  Finally, we give a recommended procedure for the derivation of abundances in the case that only spectra of limited wavelength coverage are available so that the optimal method can no longer be used. ",
        "introduction": "Powerful constraints on theories of galactic chemical evolution and on the star formation histories of galaxies can be derived from the accurate determination of chemical abundances either in individual star forming regions distributed across galaxies or through the comparison of abundances between galaxies.  The use of optical emission to estimate abundances in extragalactic ionized hydrogen (\\HII) regions dates back as early as the 1940s \\citep{Aller42}. The advent of linear detectors with high dynamic range, and the development of quantitative photoionization models provided the means whereby chemical abundances could be measured quantitatively in both galactic and extragalactic \\HII\\ regions \\citep[eg.,][]{Peimbert75,Pagel86,Osterbrock89,Shields90,Aller90}. In such work, measurements of the Hydrogen and Helium recombination lines are used along with collisionally excited lines observed in one or more ionization states of heavy elements.  Oxygen is commonly used as the reference element because it is relatively abundant, it emits strong lines in the optical regime, it is observed in several ionization states, and line ratios of frequently-observed lines can provide good temperature and density diagnostics (ie. \\ion{O}{1}~$\\lambda$6300, \\OII~$\\lambda$3727,7318,7324 \\OIII~$\\lambda$4363,4959,5007). However, the densities of extragalactic \\HII\\ regions are so low that the density-sensitive line ratios are rarely used.  Ideally, the oxygen abundance is measured directly from the ionic abundances obtained using a determination of the electron temperature of the \\HII\\ region, and including an appropriate correction factor to allow for the unseen stages of ionization. This is known as the Ionization Correction Factor (ICF) method.  Electron temperature is useful as an abundance indicator since higher chemical abundances increase nebular cooling, leading to lower \\HII\\ region temperatures.  The electron temperature can be determined from the ratio of the auroral line \\O4363\\ to a lower excitation line such as \\OIIIl. In practice, however, \\O4363\\ is often very weak, even in metal-poor environments, and cannot be observed at all in higher metallicity galaxies. In addition, \\O4363\\ may be subject to systematic errors when using global spectra.  For example, \\citet{Kobulnicky99b} found that for low metallicity galaxies, the \\O4363\\ diagnostic systematically underestimated the global oxygen abundance, while for more massive, metal-rich galaxies, empirical calibrations using strong emission-line ratios can serve as reliable indicators of the overall oxygen abundance in \\HII\\ regions.  With the rapidly increasing interest in the star formation history of the high-redshift universe, it has become much more important to determine the chemical evolutionary state of distant galaxies, even when only  spatially unresolved (global) emission-line spectra are available \\citep{Steidel96, Kobulnicky99a, Kobulnicky99b}. However given the range of temperatures, ionization parameters, and metallicities encountered in different galactic systems, abundances derived using global spectra may often be subject to important systematic errors.  For such galaxies and other star-forming regions without a measurable [{\\ion{O}{3}}] $\\lambda$4363 flux, abundance determinations become entirely dependent on the measurement of the ratios of the stronger emission-lines. These still have the potential to deliver reasonably accurate estimates. The most commonly used such ratio is \\ratioR23 (otherwise known as \\R23), first proposed by \\citet{Pagel79}. The logic for the use of this ratio is that it provides an estimate of the total cooling due to oxygen, which, given that oxygen is one of the principle nebular coolants, should in turn be sensitive to the oxygen abundance. However, the calibration of this ratio in terms of the abundance has proved to be rather difficult due to the lack of good-quality independent data.  As discussed above, \\O4363\\ is usually well-observed only for metal-poor (Z$<$0.5 \\Zsun) star-forming regions.  At higher metallicities, detailed theoretical model fits to the data must be used.  As a result, many calibrations of \\R23 are available, including \\citet{Pagel79,Pagel80,Edmunds84,McCall85, Dopita86,Torres-Peimbert89,Skillman89,McGaugh91,Zaritsky94,Pilyugin00}, Charlot01. Recent calibrations of \\R23 produce oxygen abundances which are comparable in accuracy to direct methods relying on the measurement of nebular temperature, at least in the cases where these direct methods are available for comparison \\citep{McGaugh91}. Rather than calculating model fits to the data, some authors have favoured a more empirical approach.  For example, \\citet{Thurston96} estimated the electron temperature from the easier-to-observe \\NII\\ line, which was then calibrated against \\R23. However, this has the drawback that the strength of the \\NII\\ must be calibrated separately, and we are dealing here with an element which has both primary and secondary nucleosynthesis components whose relative contributions depend on metallicity.   One drawback of using \\R23 (and many of the other emission-line abundance diagnostics) is that it depends also on the ionization parameter ($q$) defined here as: \\begin{equation} q=\\frac{S_{{\\rm H}^{0}}}{n}  \\label{1} \\end{equation} where $S_{{\\rm H}^{0}}$ is the ionizing photon flux through a unit area, and $n$ is the local number density of hydrogen atoms. The ionization parameter $q$, can be physically interpreted as the maximum velocity of an ionization front that can be driven by the local radiation field. This local ionization parameter can be made dimensionless by dividing by the speed of light to give the more commonly used ionization parameter; ${\\cal U}\\equiv q/c.$ Some calibrations have attempted to take this into account \\citep[eg.][]{McGaugh91}, but others do not \\citep[eg.][]{Zaritsky94}. Another difficulty in the use of \\R23 and many other emission-line abundance diagnostics is that they are double valued in terms of the abundance. This is because at low abundance the intensity of the forbidden lines scales roughly with the chemical abundance while at high abundance the nebular cooling is dominated by the infrared fine structure lines and the electron temperature becomes too low to collisionally excite the optical forbidden lines. When only double-valued diagnostics are available, an iterative approach which explicitly solves for the ionization parameter also helps to resolve the abundance ambiguities, as will be described in this paper.  A combination of an extended IUE database, IRAS data, and the existence of better stellar evolutionary tracks which include mass loss and overshooting allowed large advances in the modelling of starburst emission spectra in recent years \\citep{Mas-Hesse91}.  With recent advances in nebular physics, detailed photoionization models have been produced which include self-consistent treatment of nebular and dust physics \\citep[eg.,][]{Dopita96,Ferland98}.  These can be used in conjunction with modern stellar population synthesis models to synthesize the spectra of starburst galaxies from the UV to the X-ray regime.  These models have been used by us \\citep{Dopita00a,Kewley01} to simulate the emission-line spectra of \\HII\\ regions and starburst galaxies, respectively. Here, these models are used to develop an optimal scheme for abundance determination based on the range of possible combinations of bright optical or IR emission-lines which are likely to be available to the observer. ",
        "conclusions": "\\label{conclusion}  We have presented a range of diagnostics using current theoretical models for determining the abundances and ionization parameters of star forming regions.  The appropriate diagnostics to be used depends on the wavelengths observed, and therefore on the availability of particular emission line ratios.  Our diagnostics have been compared with those of \\citet{McGaugh91}, \\citet{Zaritsky94} and \\citet{Charlot01}, and we arrive at the following conclusions;  \\begin{enumerate} \\item Our \\NIIOII\\ diagnostic is clearly the best diagnostic to use in the Z$>0.5$~\\Zsun\\ (\\OH~$>8.6$) regime, as it produces a remarkably tight correlation with the abundance determined using the average of five previous techniques, none of which can reproduce such a tight correlation when used individually. The \\NIIOII\\ diagnostic produces this tight correlation as a result of both its independance of ionization parameter and its strong metallicity sensitivity.  We have investigated the effect of poorly reddening-corrected or un-corrected spectra on the derived abundance using this diagnostic.  We show that this adds both a large degree of uncertainty in the abundance derived, and a systematic bias to estimate higher than average abundances. However, we have shown that `classic' extinction correction methods such as those based on the Whitford reddening curve, provide sufficiently good extinction correction to allow the \\NIIOII\\ diagnostic to be used for reliable abundance determinations.  \\item As has been shown previously, the common abundance diagnostic \\R23\\ depends strongly on ionization parameter, and the common ionization parameter diagnostic \\OIIIOII\\ depends strongly on abundance.  An iterative approach is often required to resolve these dependancies.  Unlike previous methods, we provide techniques for explicitly determining the ionization parameter, rather than including this as a `correction factor' to the abundance diagnostic.  \\item Due to the local maximum in \\R23, for objects with abundances between $8.4 <$~\\OH~$< 8.8$ (ie $0.6 <$~log(\\R23)~$< 1.0$), a different diagnostic to ours in this range should be used to obtain a more reliable abundance estimate.  For \\OH $\\ge 8.9$, our \\R23\\ method delivers a slightly higher abundance than the comparison abundance, but this shift is equal to if not less than that seen for the other two \\R23\\ -based diagnostics (M91 and Z94) in this region. For \\OH$ \\le 8.5$, our \\R23\\ method is much more reliable than the other two \\R23\\ based methods. The C01 \\NIISII\\ method is also reliable in this region, with a similar degree of scatter.  \\item We have presented a new combined diagnostic based only on three \\R23\\ diagnostics: ours, M91 and Z94.  This diagnostic eliminates the systematic shift inherent to all three techniques used on their own and significantly reduces the scatter or rms error.  \\item The ionization parameter diagnostic [\\ion{S}{3}]/[\\ion{S}{2}] is independant of abundance, enabling a non-iterative approach to be used if \\SIII\\ and \\SII\\ are available in the spectrum. However, our current models do not allow for reliable ionization parameter diagnostics using [\\ion{S}{3}]/[\\ion{S}{2}], or abundance determinations derived from \\dS23.  We believe this is a result of either the use of an incorrect sulfur to oxygen abundance ratio, errors in the sulfur depletion factor used, or errors in fundamental atomic data.  If the sulfur lines are the only strong lines available, we recommend an empirical method such as that given in \\citet{Diaz00} or the modified Diaz \\& P\\'{e}rez-Montero diagnostic presented here which significantly reduces the errors inherent in the Diaz \\& P\\'{e}rez-Montero diagnostic alone.  \\item Both the \\NIISII\\ and \\NIIHa\\ ratios strongly depend on ionization parameter, so an ionization parameter diagnostic should be used to aid abundance determinations when using these ratios.  Neither diagnostic is very sensitive to abundance and should be used with caution. In particular, \\NIISII\\ based diagnostics have a larger degree of uncertainty than schemes based on other ratios such as \\NIIOII\\ or \\R23.  \\item Finally, for spectra of \\HII\\ regions or star-forming galaxies in which the \\NII, \\OII, \\OIII, \\SII\\, and the Balmer lines of Hydrogen are available, we present a method using a combination of techniques for optimally determining the abundance from these strong lines alone.  This method takes advantage of the reliability of our \\NIIOII\\ for the intermediate to high metallicity range, and our \\R23\\ diagnostic combined with the \\citet{Zaritsky94}, \\citet{McGaugh91} and \\citet{Charlot01} diagnostics for the lower abundance range.  This technique can be used over the entire metallicity range, and appears to be prone to much smaller intrinsic errors than all other techniques presented, including those of \\citet{McGaugh91,Zaritsky94,Charlot01}.  In addition, there is no systematic offset in the derived abundance when compared with the average of three previous techniques.  We strongly recommend use of this technique if all the required emission-lines are available, and fluxes can be extinction corrected using classical methods.   \\end{enumerate}"
    },
    "0206/astro-ph0206176_arXiv.txt": {
        "abstract": "The Cold Dark Matter (CDM) model for galaxy formation predicts that a significant fraction of mass in the dark matter haloes that surround $L\\sim L_*$ galaxies is bound in substructures of mass $10^4$--$10^7\\,\\msun$.  The number of observable baryonic substructures (such as dwarf galaxies and globular clusters) falls short of these predictions by at least an order of magnitude.  We present a method for searching for substructure in the haloes of gravitational lenses that produce multiple images of QSOs, such as 4-image Einstein Cross lenses.  Current methods based on broadband flux ratios cannot cleanly distinguish between substructure, differential extinction, microlensing and, most importantly, ambiguities in the host lens model.  These difficulties may be overcome by utilizing the prediction that when substructure is present, the magnification will be a function of source size.  QSO broad line and narrow line emission regions are approximately $\\sim 1$\\,pc and $>100$\\,pc in size, respectively.  The radio emission region is typically intermediate to these and the continuum emission region is much smaller.  When narrow line region (NLR) features are used as a normalisation, the relative intensity and equivalent width of broad line region (BLR) features will respectively reflect substructure-lensing and microlensing effects.  Spectroscopic observations of just a few image pairs would probably be able to cleanly extract the desired substructure signature and distinguish it from microlensing -- depending on the {\\em actual} level of projected mass in substructure.  In the rest-optical, the \\hb/\\oiii\\, region is ideal, since the narrow wavelength range also largely eliminates differential reddening problems.  In the rest-ultraviolet, the region longward and including Ly$\\alpha$ may also work.  Simulations of Q2237+030 are done as an example to determine the level of substructure that is detectable in this way, and possible systematic difficulties are discussed.  This is an ideal experiment to be carried out with near-infrared integral field unit spectrographs on 8-m class telescopes, and will provide a fundamentally new probe of the internal structure of dark matter haloes. ",
        "introduction": "}  Galaxy formation models cast in a $\\Lambda$CDM universe have experienced considerable success to date.  With a fixed choice of cosmological parameters, it is possible to make specific predictions, for example based on N-body simulations, regarding the structure of individual dark matter haloes.  One of the challenges in this field is to calculate the level of the continuously-infalling substructure that will survive within large haloes.  Present N-body simulations indicate that there should be plentiful substructure in a $L_*$ galaxy's halo, with masses above $10^7\\msun$; below this scale limitations in resolution becomes important. If these subhaloes contain stars, there should be significantly more dwarf galaxies near the Galaxy than are seen \\citep{moore:99, klypin:01}.  Various solutions to this excess problem have been proposed, including warm dark matter (which erases substructure on small scales; e.g. \\citealt{avila:01}); self-interacting dark matter (which causes subhaloes to evaporate; Spergel \\& Steinhardt 2000); and radiation-driven ionization of the intergalactic medium such that star formation in dwarf-size galaxies is suppressed at early epochs, as discussed in \\citet{somerville:02}, \\citet*{bullock:01}, and \\citet{klypin:99}.  If the substructure does exist, its gravitational influence should be detectable, and therefore this prediction of $\\Lambda$CDM in particular may be tested.  There is already some evidence in support of the need for substructure to explain the relative magnifications in strong lenses; disentangling its signature from other (equally interesting) effects is the scope of this paper.  From the positions of the images and lens galaxy of a 4-image lens system, a parametric smooth mass model for the lensing galaxy and halo can be computed (including external shear), along with predictions for the three independent flux ratios.  Measuring deviations between the predictions and the (differential-reddening-corrected) observed integrated fluxes in real Einstein Cross lenses, provides a direct estimate of the surface density in substructure \\citep*{metcalf:01,chiba:02,metcalf:02,dalal:02}.  Substructure has long been considered a possible explanation for inconsistencies in the observed flux ratio of B1422$+$231 \\citep*{mao:98}.  However, this approach has a vital flaw.  The predicted flux ratios can be strongly dependent on the parametric lens models used, making any measurement of the mass fraction in substructure \\cite[e.g.][]{dalal:02} suspect. The observational technique proposed here largely avoids this important problem, as well as the problems of intrinsic variability (on the timescale of the typical time-delay between images), and the effects of microlensing, with the possible change of continuum slope as well as intensity.  In \\S\\ref{sec:models} the main features of substructure and its lensing effects are summarized.  Simulations demonstrating the importance of source-size on magnification are presented in \\S\\ref{sec:lensing-simulations}, for a specific lens, Q2237+030.  The spectroscopic approach proposed here is discussed in \\S\\ref{sec:spec}, with the most difficult problems and their possible solutions addressed.  In \\S\\ref{sec:disc} we discuss selection criteria for lens candidates for this experiment, as well as modes of observations that would be sufficient for obtaining results.  The main conclusions are summarized in \\S\\ref{sec:conc}.  In the Appendix an estimate of the observational uncertainties is calculated.  All cosmological calculations are carried out for a $\\Lambda$CDM cosmology, with $\\Omega_{\\Lambda}=0.7$, $\\Omega_0=0.3$, though none of the arguments (except the level of substructure one might expect) depend strongly on this choice.  The Hubble parameter is $\\Ho= 65\\,h_{65}\\kms\\mpc^{-1}$. ",
        "conclusions": "\\label{sec:conc}  In $\\Lambda$CDM, a significant fraction of the dark matter in $L\\sim\\,L_*$ galaxy haloes is in coherent substructure, with masses of $10^4$--$10^8$\\,M$_{\\odot}$, which does not seem to have baryonic, or at any rate luminous, counterparts.  The substructure and small-scale intergalactic structure should be detectable by its gravitational signature in multiply-imaged QSOs.  If such substructure is sufficiently abundant along the images' lines of sight, the light from different emission regions of a QSO will be magnified differently. This would be a model independent signature of small scale substructure.  Broadband photometry is primarily a measurement of the continuum flux, which is affected by rapid fluctuations intrinsic to the QSO, and by microlensing caused by stars.  The magnitudes of both of these effects are at least comparably to what substructure is expected to produce, and conventionally can only be tracked by extensive monitoring.  We propose that these fundamental limitations can be largely overcome by obtaining high signal to noise spectra ($SNR\\gsim 100$) of two or all lensed images simultaneously, that include prominent NLR and BLR features.  The NLR features are used to normalise the spectra to each other, and then the relative fluxes of the BLR line(s) should be primarily due to substructure differential magnification.  Intrinsic variability in BLR lines is observed to amount to $\\sim10$--$35$\\%, but relatively smoothly, and on characteristic timescales of greater than six months in the rest frame (and 1+$z_s$ times longer in our observed frame).  Compared to the time delay of only days or weeks between the images in typical Einstein Crosses, the {\\em intrinsic} BLR fluxes can be assumed to be constant.  The ideal setup is with an integral field unit (IFU) in the near-infrared (NIR), targeting the rest-optical \\oiii\\ \\& \\hb\\ emission lines.  With the IFU setup, it is possible to get high-quality spectra of all images (and possibly of the lens, as well), simultaneously.  The alternative of a set of longslit observations can be made more efficient by aligning on pairs of lens images, positioning the slit close to the parallactic angle to minimize the effects of differential refraction.  Efficient implementation will require an 8-m class telescope, with excellent seeing conditions ($\\sim0.6''$), perhaps with adaptive optics.  This experiment will provide a concrete measurement, or at least place a severe upper limit on the fraction of dark-matter substructure in galaxy haloes and intergalactic space as predicted by the popular CDM models of structure formation.  These models are presently in a state of crisis because of the lack of observations of this substructure."
    },
    "0206/astro-ph0206389_arXiv.txt": {
        "abstract": "The first stars hold intrinsic interest for their uniqueness and for their potentially important contributions to galaxy formation, chemical enrichment, and feedback on the intergalactic medium (IGM). Although the sources of cosmological reionization are unknown at present, the declining population of large bright quasars at redshifts $z > 3$ implies that stars are the leading candidates for the sources that reionized the hydrogen in the IGM by $z \\sim 6$. The metal-free composition of the first stars restricts the stellar energy source to proton-proton burning rather than the more efficient CNO cycle. Consequently they are hotter, smaller, and have harder spectra than their present-day counterparts of finite metallicity. We present new results from a continuing study of metal-free stars from a cosmological point of view. We have calculated evolving spectra of Pop III clusters, derived from a grid of zero-metallicity stellar evolutionary tracks. We find that H-ionizing photon production from metal-free stellar clusters takes twice as long as that of Pop II to decline to 1/10 its peak value. In addition, metal-free stars produce substantially more photons than Pop II in the He~II ($E > 4$ Ryd) continuum. We suggest that large Ly$\\alpha$ equivalent widths ($W_{\\rm Ly\\alpha} > 400$ \\AA) may provide a means of detecting metal-free stellar populations at high redshift, and that He~II recombination lines ($\\lambda1640$, $\\lambda4686$) may confirm identifications of Population III. While Pop III clusters are intrinsically bluer than their Pop II counterparts, nebular continuum emission makes up this difference and may confuse attempts to discern Pop III stars with broadband colors. In a companion paper, we explore the consequences of evolving spectra of Pop III for the reionization of the IGM in both H and He. ",
        "introduction": "The first stars and their role in galaxy formation and the evolution of the intergalactic medium (IGM) define one of the frontiers of modern cosmology. Yet we do not know when the first stars formed, where they now reside, what happened to them over time, or even where to look for them. The Big Bang paradigm requires that the first stars were composed of H and He with only trace light elements and negligible carbon (Schramm \\& Turner 1998). Metal-free stars have unique properties owing to this primordial composition: they are hotter and smaller than their metal-enriched counterparts (Tumlinson \\& Shull 2000, hereafter TS). These features have interesting implications for their formation, their environment, and particularly for their spectra. Some key open issues are the formation epoch of the first stars (Ricotti, Gnedin, \\& Shull 2002), the mass distribution for stars formed from primeval gas (Bromm, Coppi, \\& Larson 2001), and their contribution to the reionization of H and He in the IGM. The recent suggestions by Lanzetta et al. (2002) that star formation may become more widespread in the distant past, the discovery of a gravitationally-lensed galaxy at $z = 6.56$ (Hu et al. 2002), and the apparent completion of H reionization by $z \\simeq 6$ (Becker et al. 2001; Djorgovski et al. 2001) and He at $z \\simeq 3$ (Kriss et al. 2001) make detailed examination of the first stars particularly timely.  With our present tools we can see back to $z \\sim 6$, roughly 1 Gyr after the Big Bang (a universe with $\\Omega_m$ =  0.3, $\\Omega_\\Lambda$ = 0.7, and $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$ is assumed here). Imaging and spectroscopic surveys that sample wide swaths of sky for telltale emission lines (Ly$\\alpha$ emitters; Hu et al.~1999; Rhoads et al.~2000) or peculiar broadband colors (Lyman-break galaxies, Steidel et al. 1999) have discovered more than 1000 galaxies at $z \\sim 3$. Searches such as the Sloan Digital Sky Survey (SDSS) are expanding the list of known QSOs at $z \\gtrsim 5$ and enabling high-redshift probes of the IGM (Fan et al.~2001; Schneider et al.~2002). Ground-based spectra can determine the star formation rate, metallicity, and mass of galaxies at $z = 3$ (Pettini et al.~2002). The first stars are believed to have formed between $z = 20 - 30$ (Gnedin \\& Ostriker 1997; Ricotti et al. 2002), while galaxies have been discovered up to $z \\sim 6$. The next observational frontier is the direct observation of the first stellar generation.  This work is part of an ongoing effort to examine the first stars from a unified stellar and cosmological point of view. We began with zero-age models (TS) of metal-free stars that revealed their potential importance to H and He reionization in the IGM. We followed that with a detailed calculation of the He~II recombination line fluxes from Population III clusters, assessing the prospects for discovering this emission as a distinctive signature of the first stars (Tumlinson, Giroux, \\& Shull 2001, hereafter TGS). We are concerned here with the properties of the first stars that affect their observational signatures and feedback on the IGM. These effects depend on the answer to a fundamental question about the first stars: what is the time evolution of their structure and radiation? We update the work of TS to compute new evolving spectra of metal-free stars and to assess the effects of evolution on the emergent spectra of Population III clusters.  In this paper, we use the terms ``Population III'' and ``metal-free'' interchangeably. We also note that there are two distinct usages of the term ``Pop III'' that have appeared in the literature. One usage connotes small protogalactic clusters that have $M_{\\rm DM} \\simeq 10^{6-7}$ \\Msun\\ and that must cool with molecular hydrogen lines in the absence of metals. We adhere to the stellar definition, in which ``Pop III'' and metal-free stars are identical regardless of location.  In \\S~2 we discuss the major components of our stellar evolution calculations. In \\S~3 we present the stellar evolution tracks that enable our evolving spectra and observational predictions. In \\S~4 we describe the modeling of the stellar atmospheres and construct evolving spectra for synthetic Pop III clusters. In \\S~5 we make detailed predictions of the observational signatures of metal-free stars and describe the realistic prospects for their discovery. In \\S~6 we draw some general conclusions. Paper II (Venkatesan, Tumlinson, \\& Shull 2002) presents the results of a cosmological reionization model that incorporates Pop III spectra as ionizing sources. ",
        "conclusions": "We have presented evolving spectra of metal-free stellar populations, based on newly calculated evolutionary tracks. We find that the evolution of Pop III stars follows the general patterns obeyed at higher metallicity, but with an overall shift to higher temperature. This gain in core and surface temperatures is primarily a result of the restricted abundance of $^{12}$C in primordial stars. We have used a grid of non-LTE model atmospheres to produce evolving spectra of synthetic Pop III clusters. These models have been used to compute the broadband colors and emission-line signatures from metal-free stellar clusters at high redshift.  Our specific conclusions fall into two categories. First, the unique composition of the first stars has the following important effects on their spectra: \\begin{itemize} \\item They produce 60\\% more H~I ionizing photons than their Pop II counterparts. \\item They can produce up to $10^5$ times more He~II ionizing photons than Pop II, which may lead to unusual radiative feedback effects on the IGM. \\end{itemize}  We wish to reemphasize here the uncertainty associated with stellar mass loss in evaluating the latter conclusion. If the ``second generation'' of stars has $Z \\sim 0.0001 - 0.001$ and perhaps significant populations of WR stars, the gain in He II ionizing photon production for a continuous star formation model including WR stars is reduced to about a factor of 10 (see Figure 6). However, since the importance of WR stars at extremely low metallicities is in doubt, we conclude that Pop III stars will be far more efficient than their Pop II counterparts at ionizing He II. The consequences of our models for reionization are explored more fully in Paper II, which uses these model spectra and a semi-analytic reionization model to assess the importance of the first stars for full or partial reionization of H~I and He~II in the IGM.  The peculiar properties of metal-free stars entail unusual observational signatures: \\begin{itemize} \\item They are expected to have unusually high equivalent widths of Ly$\\alpha$ and the He~II recombination lines, the most distinctive signature of Pop III stars and their nebular emission. These lines will probably be the primary means of detecting and identifying the first stars. \\item While their intrinsic spectra are significantly bluer than their Pop II counterparts, their broadband colors are similar, owing to the reprocessing of Lyman-continuum photons into continuous nebular emission. \\end{itemize}  Recent Ly$\\alpha$ surveys suggest that metal-free stars may have already been found. In future ground- and space-based searches beyond the optical bands, Ly$\\alpha$ emission will probably serve as the ``signpost'' of metal-free galaxies. Follow-up detections of the He~II lines could confirm this result. The substantial degeneracy in the broadband colors will probably leave emission line techniques as the best way of distinguishing the first stars."
    },
    "0206/astro-ph0206340_arXiv.txt": {
        "abstract": "The usage of a large amount of CdTe(CdZnTe) semiconductor detectors for solar neutrino spectroscopy in the low energy region is investigated. Several different coincidence signals can be formed on five different isotopes to measure the \\bes neutrino line at 862 keV in real-time. The most promising one is the usage of \\cd resulting in 89 SNU. The presence of \\tehfz permits even the real-time detection of pp-neutrinos. A possible antineutrino flux above 713 keV might be detected by capture on \\cdhs . ",
        "introduction": "Over the last years striking evidence arose for a non-vanishing neutrino rest mass (for reviews see \\cite{zub98,samoil}). They all come from neutrino oscillations experiments. Among them is the long standing evidence of a solar neutrino deficit also being of fundamental importance for stellar astrophysics. The deficit is seen in radiochemical detectors, namely GALLEX/GNO \\cite{gallex,kirsten} and SAGE \\cite{gavrin} using \\gaes , still the only pp-neutrino detectors available, and the Homestake experiment using \\clsd \\cite{ray}. A reduced $^8$B \\nel flux is measured by two water Cerenkov detectors, namely Super-Kamiokande \\cite{smy} and SNO \\cite{snocc}. A difference in measured fluxes among the latter resulted in evidence for an active neutrino flavour coming from the sun besides \\nel. This is due to the fact that Super-Kamiokande is using neutrino-electron scattering and SNO inverse \\bdec{} for detection. Recent SNO results of neutral current reactions on deuterium dramatically confirm the existence of further active neutrinos, being the dominant solar neutrino flux \\cite{snonc}. The solution of the solar neutrino deficit has to come from particle physics, the scenario discussed most often is neutrino oscillations. Taking all experimental results together the largest effects in the solar neutrino spectrum implied by the various oscillations solutions show up in the region below 1 MeV. Furthermore this region corresponds to 99 \\% of the solar neutrino flux and is still very important for understanding stellar energy generation \\cite{bah02}. An interesting idea to measure such low energy neutrinos in real time is suggested by \\cite{rajulens} using coincidence techniques for neutrino capture on nuclei. It is already finding its practical application in the LENS \\cite{lens}, SIREN \\cite{siren} and MOON projects \\cite{moon}. aiming to measure pp-neutrinos in real time. The technique relies on using either a large amount of double beta isotopes (\\yhss in case of LENS, \\gdhs in case of SIREN and \\moeh in case of MOON) or highly forbidden beta decay emitters like \\inhf (4-fold forbidden, currently under study in LENS as an alternative to \\yhss ), as target material \\cite{rajulensin}. Clearly an interesting spin off is the investigation of double beta decay \\cite{zub00}.\\\\ In this paper the possibility to apply the same technique for CdTe(CdZnTe) semiconductor detectors and their feasibility for solar neutrino detection is explored. CdTe semiconductor detectors have already a wide field of application in $\\gamma$-ray astronomy and medical physics. The study performed here is motivated by the COBRA project \\cite{cobra}, planning to use large amounts of CdTe-detectors for double beta decay searches. The usage of large amounts of semiconductors for solar neutrino detection was also considered in the past for Ge-detectors \\cite{laura} and GaAs \\cite{bowles} relying largely on the detection of electrons from neutrino-electron scattering. In the case discussed here, we focus on the detection of \\bes and pp-neutrinos only. Measurements of higher energetic neutrino flux components are also possible but will not be discussed. Also a possible real-time detection of pp-neutrinos via neutrino-electron scattering as well as contributions from the CNO cycle are not considered. ",
        "conclusions": "The prospects of various isotopes of Cd, Zn and Te for low energy solar neutrino spectroscopy are explored. To obtain a reasonable signal various coincidence tags can be used, as compiled in Tab.~\\ref{tab:comparison}. It allows the detection of $^7$Be in real time for five isotopes and therefore offers redundancy in the obtained results. The most promising detection signal is the ground state transition of \\cd to \\inhs resulting in 89 SNU. This has to be seen as a lower limit because CNO contributions are not taken into account. In addition $^{125}$Te allows a real time detection of pp-neutrinos with a threshold of 330 keV. The usage of semiconductors is advantageous for background reduction for $^7$Be detection is  because the monoenergetic electron forming the first step of the coincidence can be measured with good precision. Rates for excited state transitions cannnot be determined reliably because a lack of knowledge in the corresponding GT matrix elements, a problem also known from other low energy solar neutrino experiments. It might be worthwile to consider an experimental program to measure these matrix elements, which would also be valueable for double beta decay. As common for solar neutrino detection detector sizes of tons have to be considered, this kind of experiment is not feasable in the very near future."
    },
    "0206/astro-ph0206030_arXiv.txt": {
        "abstract": "{We construct and investigate the pulsar luminosity function using the new catalogue which includes data for 1315 radio pulsars. The luminosity functions are constructed for 400 and 1400 MHz separately, and they are compared. Also, the luminosity functions excluding the binary millisecond pulsars and other pulsars with low magnetic fields are constructed. The 1400 MHz luminosities as a function of characteristic age and as a function of magnetic field for radio pulsars, anomalous X-ray pulsars and dim radio quiet neutron stars are presented and the implications of the pulsar luminosity function on these new kind of neutron stars are discussed. ",
        "introduction": "As known, radio pulsars (PSRs) do not radiate isotropically, but radiate in a beam. The beaming angle in the radio band is much less than the ones in the X-ray band (Lyne \\& Graham-Smith 1998). Now, we have the data for 1315 PSRs (Guseinov et al. 2002a). From these PSRs, about 30 of them with characteristic ages $\\tau <$10$^7$ yr are observed also in optical, X-ray and $\\gamma$-ray bands. Becker \\& Tr\\\"{u}mper (1997) gives the list for 19 radio PSRs with ages $\\tau <$10$^7$ yr observed also in the X-ray band with the $ROSAT$ satellite.  For the last 10 years, very important objects in astrophysics; so called soft gamma repeaters (SGRs), anomalous X-ray pulsars (AXPs) and dim radio quiet neutron stars (DRQNSs) has been investigated very intensively (Mereghetti et al. 1996; Kaspi et al. 1996; Thompson \\& Duncan 1995; Brazier \\& Johnston 1999; Halpern et al. 2002). There is no radio radiation observed from these sources (Gaensler et al. 2001).  Some of these single neutron stars, namely, SGRs and AXPs have X-ray spin periods between 5-10 s and their magnetic field strengths are believed to be B=10$^{14}$-10$^{15}$ G (Thompson \\& Duncan, 1995). It is believed that the age of AXPs and SGRs absolutely are not more than 10$^5$ yr (Tagieva \\& Ankay, 2002). Since PSRs with characteristic time $\\tau < 10^5$ yr have magnetic fields mostly of 10$^{12}$-10$^{13}$ G and since no PSR has a magnetic field as high as 10$^{14}$-10$^{15}$ G, AXPs and SGRs are named as magnetars. It is important to note that now the properties of AXPs and SGRs are mostly investigated in the frame of the magnetar model, however, we do not know definitely that these neutron stars have such huge magnetic fields. We wonder whether the non-detection of radio radiation from SGRs, AXPs and DRQNSs with spin periods P$\\sim$5-10 s is due to the strong magnetic fields?  There are about 10 dim X-ray point sources. Some of these sources with small periods are connected to supernova remnants (SNRs) (see Gaensler et al. 2001 and Tagieva \\& Ankay 2002 for review). There is no radio radiation detected from these sources in spite of their small distances from the Sun. We may expect that the reason that there is no radio radiation from DRQNSs with small periods is because of the radio luminosity function.  For most of the SNRs closer than 5 kpc, there have been searches for pulsars (Gorham et al. 1996, Lorimer et al. 1998). However, it has been very hard to detect them. Up to 5 kpc from the Sun, only 6-7 pulsars are found to be genetically connected to SNRs (Kaspi \\& Helfand 2002, Camilo et al. 2002, Guseinov et al. 2002b). Most of them are discovered in surveys. There are about 79 SNRs up to the distance 5 kpc (Guseinov et al. 2002b).   This shows that on the average, for SNRs closer than 5 kpc, only 1 genetic connection exists for 11-13 SNRs. Is not this a result of the luminosity function and background radiation? In this paper, we aim at finding the improved luminosity function at 400 MHz and for the first time at 1400 MHz which is more important today due to the fact that the radio observations of AXPs, SGRs and DRQNSs are conducted at 1400 MHz. We also aim at investigating the properties of the luminosity function for all PSRs and single born PSRs, separately. ",
        "conclusions": "In Sections 2 and 3, the luminosity function is found for PSRs with luminosity known at 400 and 1400 MHz. Luminosity functions are found both for all PSRs and for PSRs which are born single and have characteristic ages $<$10$^7$ yr. Before, the luminosity function for 1400 MHz could not be constructed because the number of PSRs observed at this frequency was few. Now since the number of such PSRs are 862 in the Galaxy (Guseinov et al. 2002a), we constructed the luminosity function at 1400 MHz. To compare the luminosity functions in Figures 3-6, we brought them together in Figure 7. As known, using the luminosity function, the ratio of space density of PSRs with different luminosities can be determined. To find the real density of PSRs in any luminosity interval or PSRs with luminosities higher than any given luminosity, the function must be calibrated by considering the number density of PSRs which have luminosity higher than a chosen luminosity value. In Figure 7, we calibrated the luminosity functions such that they have the same values at Log L=-1. Using this function we may have information about the relative number of PSRs with different values of luminosity values.   As seen from Figure 7, the luminosity function for single born PSRs (lines 5 and 6) and luminosity function for all PSRs (lines 3 and 4) are similar at the same frequencies. There is huge difference when we compare the luminosity functions at 400 and 1400 MHz. The reason of this huge difference is that the spectral index of PSRs is sharp. By comparing the lines 3 and 5 in the figure, we see that the number density of PSRs which are born single with Log L$_{400}$ close to 1.5 is less but number density of single PSRs with Log L$_{400} \\sim$ 0-0.5 is more. On the other hand, if we compare the lines 4 and 6 we see that single born PSRs with Log L$_{1400} >$0.5 is more. There might be some doubts in the comparison of the line 3 and 5 which are closer to each other, however, the results obtained by comparison of the line 4 and 6 are must be trustworthy.  If the AXPs and SGRs are magnetars then the non-detection of radio radiation from them must be explained. If AXPs are accreting neutron stars, then they will not emit any radio radiation. The upper limits of radio fluxes at 1400 MHz for AXPs are known (Gaensler et al. 2001). In Figure 8 and 9, Log L$_{1400}$ values versus Log $\\tau$ and Log B for PSRs with ages $<$3 10$^5$ yr are represented. As seen from these figures, the radio luminosities of PSRs practically do not depend on $\\tau$ or B. Upper limit values for radio luminosities of 4 AXPs on the average are smaller than that of young PSRs. AXPs 1E2259+586 and 4U0142+625 have very small upper limits for luminosity. As we see from Figure 6 number of PSRs with such small luminosity must be very small. Therefore we may believe that AXPs and may be also SGRs have very small radio luminosity.  In Figures 8 and 9 also upper limit for radio luminosities of DRQNSs are presented. As seen from figures, without any doubt, this class of neutron stars practically do not have radio radiation. The Number of PSRs younger than 10$^5$ yr up to the distance 1 kpc is 2. The number of DRQNSs with such ages and distances is also 2. This shows that the birth rate of DRQNSs in the Galaxy must be closer to the birth rate of radio PSRs. This situation does not change if we take the beaming factor into account. The beaming factor for radio PSRs is 1/2 which implies the number of radio PSRs younger than 10$^5$ yr and closer than 1 kpc to be 4. $CHANDRA$ satellite observations will considerably increase the number of known DRQNSs."
    },
    "0206/astro-ph0206206_arXiv.txt": {
        "abstract": "{A short analysis is presented of the effects on the cepheid light curve shape, i.e. on the Fourier parameters usually adopted for its description, of the blending of the stellar image with other close stars. The conclusion is that, within reasonable error, the effects are in general small and the Fourier decomposition is confirmed to be a useful tool for pulsation mode discrimination. A large effect has been found on the phase differences in a narrow period range corresponding to the known resonance centers between pulsation modes. ",
        "introduction": "Cepheids are primary distance indicators for external galaxies and those used for this application pulsate in the fundamental mode. First overtone mode Cepheids are brighter by about 0.4 mag than fundamental mode pulsators with the same period. Since the period--luminosity relation has an intrinsic dispersion, which depends on several parameters (e.g. different effective temperature or color, different reddening, contribution from stellar companions), it is essential to remove the contaminating stars that are pulsating in a different mode. The large surveys of the Magellanic Clouds performed by MACHO (e.g. Welch et al. \\cite{we}), EROS (e.g. Beaulieu et al. \\cite{bea}) and OGLE (e.g. Udalski et al. \\cite{uda4}) projects proved that the Fourier decomposition is a good technique for discriminating the mode among short period ($P \\la 6$ d) Cepheids. More recently, the technique began to be applied to Cepheids of farther galaxies in the Local Group, such as IC 1613 (e.g. Antonello et al. \\cite{pa2}; Dolphin et al. \\cite{dol}) and M33 (Mochejska et al. \\cite{moc1}).  The large surveys offered also the opportunity of discussing the problems related to blending. Mochejska et al. (\\cite{moc}) define the blending as the close projected association of a Cepheid with one or more intrinsically luminous stars, which cannot be detected within the observed point-spread function by the photometric analysis. There is some debate about the implications for the distance determination related to the blending and more generally to poor resolution of the stellar images in these galaxies. The blending also has other effects on the light and the color curves. Mochejska et al. (\\cite{moc}) note that in the case of a red or blue companion the light curve exhibits a flatter minimum. As regards binaries, it is well-known that the observed amplitude of the light curve is affected by the luminosity of a bright companion. Could it be that the blending, apart from producing a lower amplitude, also mimics a different pulsation mode? Recently, we recalled that in principle such an effect on the Fourier parameters is small in the context of mode identification (Antonello et al. \\cite{afm}). Here we report the results of simulations that support this conclusion, and we discuss some unexpected characteristics.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{figura1.ps}} \\caption[ ]{Lower panel: blending effect on the $V$ light curve of a Cepheid, for different values of the luminosity of the companion star ($L_c$). Upper panel: comparison between the light curve for $L_c=0$ (continuous line) and $L_c=4<L>$ (dotted line) scaled to the same amplitude. } \\label{fit} \\end{figure} ",
        "conclusions": "\\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{figura2.ps}} \\caption[ ]{The plots show how the Fourier parameters and light curve amplitude of a Cepheid change according to the luminosity of a companion star ($L_c/<L>$ is the ratio of the luminosity of the companion to the average value of the Cepheid). The symbols indicate different values of the mean error $\\sigma$ of measurements adopted in the simulations: {\\em filled circle:} $\\sigma$=0, {\\em filled triangle:} $\\sigma$=0.02; {\\em open circle:} $\\sigma$=0.05; {\\em open triangle:} $\\sigma$=0.1 mag. The errorbar indicates the formal error of the respective parameter. $\\Delta{mag}$ is the average magnitude difference between the Cepheid and the blended image } \\label{four} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{diffe1.eps}} \\caption[ ]{Simulated blending effect on the $I$-band light curves of all the OGLE fundamental mode Cepheids in the SMC. The plots show the difference of $R_{21}$ and $\\phi_{21}$ between the light curves for $L_c=2<L>$ and $L_c=0$ } \\label{diffe1} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{diffe2.eps}} \\caption[ ]{Simulated blending effect on the $I$-band light curves of all the OGLE first overtone mode Cepheids in the SMC. The plots show the difference of $R_{21}$ and $\\phi_{21}$ between the light curves for $L_c=2<L>$ and $L_c=0$ } \\label{diffe2} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{diffe3.eps}} \\caption[ ]{ The difference of $R_{41}$ and $\\phi_{41}$ between the simulated light curves for $L_c=2<L>$ and $L_c=0$ of fundamental mode Cepheids. It should be compared with Fig. \\ref{diffe1} } \\label{diffe3} \\end{figure} The cases discussed here concern reasonable light curves; we do not consider the problems related to very faint variables, which can hardly be detected at minimum light. The requirement is that in the $P$ interval where it is possible to find stars pulsating in different modes, the Fourier parameters must allow us to make the discrimination. It is known that this occurs for $P \\la 6$ d for the fundamental and the first overtone mode, using only light curves parameters. The results of the simulations show that in this $P$ range the blending has a negligible effect when we compare the differences introduced by it with the size of the parameters themselves. In particular, a blended fundamental mode pulsator will have slightly larger amplitude ratios than a non-blended one; we recall that the amplitude ratios of fundamental mode pulsators are intrinsically larger than those of first overtone mode ones in this $P$ range. The same occurs for a first overtone mode pulsator compared with a second overtone one, for $P \\la 1.3$ d. On the other hand, a heavily blended first overtone pulsator increases its $R_{21}$ value, but in general not so much so as to be confused with a fundamental mode pulsator. In conclusion, the blending due to various reasons is not an issue for the pulsation mode discrimination.  The color of the companion stars is not relevant for the present discussion, as long as their contribution is constant; some (second order) effects could be related to their intrinsic variability, both in terms of photometric variability and/or Doppler shift. The influence of the photometric variability of the companion itself can be usually accurately estimated, since an adequate time series analysis is sufficient to disentangle the different contributions, because of the different periodicities or timescales involved. Also in this case, however, it is wise to work with intensities rather than with magnitudes. Variable seeing conditions could have some effect on the estimate of the intensity through the PSF fitting procedure; however in this case we would expect just an increased error in the measurement.  The plots in Figs. \\ref{diffe1} and \\ref{diffe2} suggest some interesting considerations. A light curve with an altered value of the mean luminosity, such as that depicted in Fig. 1, or expressed with a different, nonlinear mathematical function (e.g. the intensity instead of the magnitude) is characterized of course by (usually slightly) different Fourier parameters. If we estimate the differences related to these changes, we note that the largest ones are for the phases of the Fourier components with smaller amplitude; for example, at about 10 d some stars have $R_{21} < R_{i1}$, for $i$ from 3 up to 6 or more. The large differences are not due to errors or to uncertainties, since here we are not dealing with observed data but with synthetic light curves (i.e. the fitting curves), which are in principle error--free. In other words the differences are {\\em intrinsically real} and reflect directly the change of the shape introduced by the different mathematical function. The interpretation of this feature is reported in the Appendix; from that, we conclude that the observed dispersion is strictly related to the smallness of the Fourier component involved. In our example, the small second Fourier component has changed its phase value by several tenths of a radian, while for the other components the change is much smaller. For the same reason we should expect an analogous results for $\\phi_{41}$, i.e. we should have some dispersion at $P \\sim 7$ d, where $R_{41}$ is small since another resonance, $P_{0}/P_{4}=3$, should be operating there (e.g. Antonello \\cite{ant}). Indeed this is shown in Fig. \\ref{diffe3}; note also that the discontinuity of $\\Delta{R_{21}}$ located at 10 d is replaced by that of $\\Delta{R_{41}}$ at about 7 d. In a certain sense, plots such as those shown in Figs. \\ref{diffe1}, \\ref{diffe2} and \\ref{diffe3} are better indicators of resonance effects than the classical ones, because they are free of subjective corrections of the phase differences by $\\pm{2\\pi}$, which could be uncertain, mainly for the higher orders. Finally, it is possible to note two minima in the lower panel of Fig. \\ref{diffe1}, one at the resonance center, and the other at $\\log P \\sim 1.5$. Kanbur et al. (\\cite{kan}) noted the structural change of the light curves at this $P$; these features still await a theoretical interpretation.  Last but not least, we remark further that several problems with the time series analysis of stellar luminosities would be simplified by adopting intensity scales instead of magnitude scales. This statement is not new, of course. Our comment is just further support to the proposal of abandoning the magnitudes. In fact, the blending has no effect on the light curve shape when we use intensity light curves, and this is an advantage, since one is always dealing with observed parameters which are affected by errors."
    },
    "0206/astro-ph0206212_arXiv.txt": {
        "abstract": "{ The spectral evolution of powerful double radio galaxies (FR II's) is thought to be determined by the acceleration of electrons at the termination shock of the jet, their transport through the bright head region into the lobes and the production of the radio emission by synchrotron radiation in the lobes. Models presented to date incorporate some of these processes in prescribing the electron distribution which enters the lobes. We have extended these models to include a description of electron acceleration at the relativistic termination shock and a selection of transport models for the head region. These are coupled to the evolution of the electron spectrum in the lobes under the influence of losses due to adiabatic expansion, by inverse Compton scattering on the cosmic background radiation and by synchrotron radiation. The evolutionary tracks predicted by this model are compared to observation using the power/source-size ({\\em P-D\\/}) diagram. We find that the simplest scenario, in which accelerated particles suffer adiabatic losses in the head region which become more severe as the source expands produces {\\em P-D\\/}-tracks which conflict with observation, because the power is predicted to decline too steeply with increasing size.  Agreement with observation can be found by assuming that adiabatic losses are compensated during transport between the termination shock and the lobe by a re-acceleration process distributed throughout the head region. ",
        "introduction": "Powerful double radio galaxies or \\lq classical doubles\\rq\\ (CDs) owe their name to the extended (hundreds of kpc) lobes of radio emission they exhibit on opposite sides of the parent galaxy. \\citet{fanaroffriley74} classified these sources as type II objects; they have luminosities $P_{178\\,{\\rm MHz}}>5 \\times 10^{25}\\,{\\rm W\\,Hz}^{-1}$ and are edge-brightened, with bright outer hotspots. It is universally agreed that the radio continuum of the CDs is synchrotron radiation from relativistic electrons and perhaps positrons. The standard scenario is that of a jet propagating from the galaxy to the outer parts of the lobes, passing through a shock front at the hotspots and subsequently filling a \\lq cocoon\\rq\\ around the jet with radiating particles --- see, for example, \\citet{begelmancioffi89,peacock99}.   \\citet{falle91} considered radio sources as expanding bubbles that are fed by supersonic jets and drive a bow shock into a radially stratified external medium. He showed that the jet length and the bow shock grow in a self-similar way for external atmospheres in which the density drops off more slowly than $1/r^2$ from the centre of the galaxy. Subsequently, \\citet{kaiseralexander97} showed that the cocoons also expand self-similarly. Based on this result, \\citet{kaiseretal97} developed an analytical model for the spectral evolution of FR~II sources as a function of redshift, jet power and the scaling of the external density profile. They assumed that electrons with an initial power-law distribution in energy are continuously injected into the plasma immediately downstream of the shock front that terminates the jet. A key point of their treatment is that the pressure in this region equals that which drives the bow shock into the external medium on the axis of the source. This results from the fact that the jet is confined by the pressure of the cocoon \\citep{begelmancioffi89}, so that the working surface over which its thrust is distributed expands self-similarly, along with the bow shock and cocoon. Particles undergo adiabatic losses on moving from the terminal shock front into the main part of the cocoon, whose pressure is lower than that at the working surface by a constant factor (taken as $\\approx16$ by \\cite{kaiseretal97}). They then undergo synchrotron and inverse Compton losses, as well as adiabatic losses as the lobe of the source expands. Using this model, \\cite{kaiseretal97} computed evolutionary tracks in the power-linear size ({\\em P-D}) plane. These display decreasing power as the size increases, in agreement with observations, which indicate a deficit of large luminous sources.  However, \\cite{blundelletal99} noted that there is no evidence that the size of the working surface is proportional to the size of the source. Instead, the hot-spots of all known FR~II sources appear to be a few kiloparsecs in diameter. This has a profound influence on the predicted evolutionary tracks, because it implies that the adiabatic losses suffered upon moving into the cocoon by particles accelerated at the termination shock are not constant, but increase strongly with the age of the source. \\citet{blundelletal99} modelled the spectral evolution of CDs including this effect by prescribing the spectrum of electrons entering the cocoon to be a broken power-law distribution. They, too, found evolutionary tracks that agree with the observed lack of large, luminous radio sources, as well as with several other properties of the samples they investigated.  In this paper we present computations of evolutionary tracks based on the picture adopted by \\citet{blundelletal99}. However, instead of prescribing the electron distribution that enters the cocoon, we assume that electrons are accelerated by the first-order Fermi process at the termination shock and then propagate through the hot spot region into the cocoon or \\lq lobe\\rq\\ according to one of two models, which we designate case A and case B. In case~A, the full adiabatic energy loss corresponding to the age dependent pressure difference between the hot spot and cocoon is applied. This results in $P$--$D$ tracks which are in conflict with the observations. One possible way out of this problem is to assume that a re-acceleration process occurs after the initial encounter with the termination shock. Such a process is indicated independently by investigations of the spectra of individual hot-spots and of optical synchrotron emitting jets \\citep{meisenheimeretal96,meisenheimeretal97,perleyetal97,wagnerkrawczynski00}. This motivates our case~B model, in which we assume that the adiabatic losses between the termination shock and the lobe are compensated by a re-acceleration process during propagation through the head region. In both cases electrons are carried through the high-loss region by fluid elements whose residence time is distributed according to a specified transport equation. During this time they suffer synchrotron and inverse Compton losses. After entering the cocoon, we follow the particle distribution as it cools and radiates and compute evolutionary tracks as well as snapshots of the spectrum at different ages.  The paper is organised as follows: in Sect.~2 we summarise the hydrodynamic picture which we adopt for the source evolution. In Sect.~3 our treatment of Fermi acceleration is described, the transport model is specified and the kinetic equation obeyed by particles in the cocoon is formulated. Section~4 is devoted to a series of tests and special cases to illustrate the way in which the transport and acceleration models influence the particle distribution in the hot spot and the lobe. Our main results on spectral evolution are presented in Sect.~5, where we conclude that the observed data exclude case~A i.e., they can be explained only by a model such as case~B, that includes re-acceleration of electrons after their encounter with the termination shock. This conclusion and our results on source spectra at different ages are discussed and compared with previous work in Sect.~6. ",
        "conclusions": "We have introduced a new description of the acceleration and transport of particles in the hot spots and lobes of FR~II radio galaxies. Using the standard approach of assuming a power-law spectrum injected at the termination shock, we model the subsequent transport through the head region into the lobes using a formalism which permits the investigation of \\lq anomalous\\rq\\ regimes. These transport regimes are described by the single parameter $\\alpha$, which determines the time-dependence of the mean-square displacement of a particle: $\\left<\\Delta r^2\\right>\\propto t^\\alpha$. Standard diffusion, in which the flux is proportional to the gradient of the particle density (Fick's law) corresponds to $\\alpha=1$. The anomalous regimes of sub-diffusion and supra-diffusion correspond respectively to $\\alpha<1$ and $\\alpha>1$. In the current application, the most important physical difference between these regimes is that they permit a range of escape times for the particles from the high-loss head region. Compared to the case of standard diffusion, these times have a much wider distribution in sub-diffusion and a much tighter one in supra-diffusion. The new formalism also requires a parameter $\\td$ that corresponds to the average residence time in units of the cooling time. Thus, in total, one more parameter is used than the standard \\citet{kardashev62} model with a single escape time for all particles and one parameter fewer is needed than the model introduced by \\citet{blundelletal99}, which requires two break frequencies and a \\lq leakiness\\rq.  The self-similar description of the hydrodynamical evolution of classical double radio sources implies that the lobe pressure falls with time. Because the pressure in the primary hot-spots does not appear to undergo this evolution, the particles injected at there in older sources have to overcome ever more severe adiabatic losses. These losses were not modelled by \\cite{kaiseretal97}, who specified the distribution of particles entering the lobes. \\citet{blundelletal99} pointed out their importance and included them in their $P$--$D$ tracks. However, because the transport model did not include an upper cut-off of injected particles, but merely a spectral break, they were still able to account for the observed population of sources. We find that the adiabatic losses between primary hot spot and lobe prevent this when our transport model is adopted. The large, high luminosity sources can be explained only as very powerful jets in very dense environments, in which case they should have smaller, brighter predecessors, that are absent from the 3CR revised catalogue. There could be several reasons for this. One possibility is that at high redshift, small sources with powerful jets exist, but are not bright in the radio because of absorption. The missing progenitors could then be powerful, high redshift analogues of the \\lq\\lq Gigahertz Peak Spectrum\\rq\\rq\\ sources \\citep{bicknelletal97}. Another is that the jet power $Q_0$ remains quite weak during the early evolution of a high-redshift source, becoming strong only after $\\sim10\\,$Myr. Intermittent jet activity in FR~II sources has already been proposed \\citep{reynoldsbegelman97}, but as a solution to the opposite problem of an observed {\\em over}-abundance of small powerful sources (at lower redshift). In general, it seems easier to imagine the jet power to be larger during earlier phases of a source's life, leading to a faster decay of the specific source power with time.  However, the solution we propose to this problem is that acceleration is not confined to the primary hot spot, but occurs throughout the head region. Independent evidence from individual sources in the infra-red, optical and X-ray bands \\citep{meisenheimeretal96,meisenheimeretal97,perleyetal97,wagnerkrawczynski00} supports this idea. A plausible mechanism is repeated encounters with weaker shock fronts which may permeate this turbulent region. We have investigated the $P$--$D$ tracks and spectra of a model in which it is assumed that this distributed acceleration compensates the expansion losses between primary hot spot and lobes, without significantly modifying the spectrum. Our tracks are more akin to those presented by \\cite{kaiseretal97} and we have investigated their sensitivity to the new transport parameter $\\alpha$. We have shown that some of the properties noted by \\citet{blundelletal99}, such as a redshift/spectral index and a size/spectra index correlation also emerge from our model."
    },
    "0206/astro-ph0206024_arXiv.txt": {
        "abstract": "The outstanding capabilities of the {\\sl Chandra} X-ray observatory have greatly increased our potential to observe and analyze thermal radiation from the surfaces of neutron stars (NSs). Such observations allow one to measure the surface temperatures and confront them with the predictions of the NS cooling models. Detection of gravitationally redshifted spectral lines can yield the NS mass-to-radius ratio. In rare cases when the distance is known, one can measure the NS radius, which is particularly important to constrain the equation of state of the superdense matter in the NS interiors. Finally, one can infer the chemical composition of the NS surface layers, which provides information about formation of NSs and their interaction with the environment. We will overview the recent {\\sl Chandra} results on the thermal radiation from various types of NSs --active pulsars, young radio-quiet neutron stars in supernova remnants, old radio-silent ``dim'' neutron stars-- and  discuss their implications. ",
        "introduction": "Observational study of thermal emission from isolated (non-accreting) neutron stars (NSs) became possible after the launch of {\\sl Einstein} (1979) and {\\sl EXOSAT} (1983), the first X-ray observatories equipped with X-ray telescopes. Contrary to nonthermal radiation of NSs, generated in the pulsar magnetospheres and observed from radio to $\\gamma$-rays, thermal emission originates immediately at the surface, with the bulk of the energy flux in the soft X-ray band. Therefore, comparing the spectra and the light curves with the models for NS thermal radiation (see the contribution by Zavlin \\& Pavlov in these Proceedings; ZP02 hereafter), one can infer the NS surface temperature, magnetic field, gravitational acceleration, and chemical composition, as well as the NS mass and radius. The analysis of such observations allows one to trace thermal evolution of NSs (Tsuruta 1998) and constrain the properties of the superdense matter in the NS interiors (D.G.\\ Yakovlev et al., these Proceedings).  Observational manifestations of NSs are very diverse, and only some of these exotic objects are suitable for observing the NS thermal emission. Most of the {\\em detected} NSs are radio pulsars. Very young, active pulsars (with ages $\\tau\\lapr 10^3$ yr --- e.g., the Crab pulsar) are quite hot, with expected surface temperatures $\\sim 1$--2 MK, but their nonthermal radiation is so bright that the thermal radiation is hardly observable. The nonthermal component in X-ray emission of middle-aged pulsars ($\\tau\\sim 10^4$--$10^6$ yr --- e.g., B0656+14, B1055--52) is much fainter, and their thermal radiation, with temperatures 0.3--1 MK, can dominate at soft X-ray and UV energies. The surfaces of old pulsars ($\\tau\\gapr 10^6$ yr --- e.\\ g., B0950+08, B1929+10, J0437--4715) are too cold to be seen in X-rays, but their polar caps are expected to be heated up to a few million kelvins, hot enough to emit observable thermal X-ray radiation.  In addition to active pulsars, a number of radio-quiet isolated NSs emitting thermal-like X-rays have been detected, with typical temperatures $\\sim 0.5$--5 MK. They are usually subdivided in four classes: Anomalous X-ray Pulsars (AXPs; see S.\\ Mereghetti et al., these Proceedings), Soft Gamma-ray Repeaters (SGRs; Kouveliotou 1995), ``dim'' or ``truly isolated'' radio-silent NSs (i.e., not associated with supernova remnants [SNRs]; Treves et al.\\ 2000), and compact central sources (CCOs) in SNRs (Pavlov et al.\\ 2002a) which have been identified with neither active pulsars nor AXPs/SGRs. Their observational manifestations (particularly, multiwavelength spectra) are quite different from those of ``active'' pulsars, and their properties have not been investigated as extensively, but the presense of the thermal component in their radiation provides a clue to understand the nature of these objects.  First thermally emitting isolated NSs were detected with {\\sl Einstein} (e.g., Fahlman \\& Gregory 1981; Cheng \\& Helfand 1983; Helfand \\& Becker 1984; Seward, Charles, \\& Smale 1986; C\\'ordova et al.\\ 1989) and investigated with {\\sl EXOSAT} (Brinkmann \\& \\\"Ogelman 1987; Kellett et al.\\ 1987). In 1990's, \\ros\\ and \\asc\\ detected X-rays from more than 30 radio pulsars, including a few thermal emitters (\\\"Ogelman 1995), and discovered several radio-silent NSs (see Becker \\& Pavlov 2002 for a review). Thermal emission from a few isolated NSs was also detected in the optical-UV range with the {\\sl Hubble Space Telescope} (e.g., Pavlov, Stringfellow \\& C\\'ordova 1996; Walter \\& Matthews 1997).  The new era in observing X-ray emission from NSs has started with the launch of two currently operating X-ray observatories of outstanding capabilities, {\\sl Chandra} and {\\sl XMM}-Newton. In this review we present recent results on thermal X-ray emission from isolated NSs of various types observed with \\cha. A more general overview of \\cha\\ observations of NSs is presented in this volume by M.C.\\ Weisskopf, while the {\\sl XMM}-Newton results on NSs are reviewed by W.\\ Becker.  \\begin{figure}[ht] \\centerline{\\psfig{file=pavlov_fig_1.ps,height=10cm,clip=} } \\caption{ \\cha\\ HRC-I image of the Vela pulsar and surrounding nebula. } \\label{velaimg} \\end{figure}  \\begin{figure} \\centerline{\\psfig{file=pavlov_fig_2.ps,height=12cm,clip=} } \\caption{ Two-component (magnetic hydrogen NS atmosphere plus power law) model fit to the combined HRC-S/LETG and ACIS count rate spectra. The bottom panel shows the contributions from the thermal and nonthermal components (lilac and green lines, respectively). } \\label{velaspec1} \\end{figure}  \\begin{figure} \\centerline{\\psfig{file=pavlov_fig_3.ps,height=7cm,clip=} } \\caption{ Multiwavelength energy spectrum of the Vela pulsar.  The solid line shows the NS atmosphere plus PL fit to the observed LETG and ACIS spectra. The dotted line corresponds to same model spectrum corrected for the interstellar absorption and extrapolated to lower and higher energies. The dash-dot lines show the extrapolated optical and EUV absorbed spectra (see Pavlov et al.~2001b for details and references). } \\label{velaspec2} \\end{figure}  \\begin{figure} \\centerline{\\psfig{file=pavlov_fig_4.ps,height=6cm,clip=} } \\caption{Light curve of the Vela pulsar from two HRC-I observations of January and February 2000. The zero phase corresponds to the radio peak. } \\label{velalc} \\end{figure} ",
        "conclusions": "Because of the limited space, we described the observational results on only a handful of observed thermally emitting isolated NSs. We only briefly mentioned AXPs and SGRs, whose X-ray radiation almost certainly contains a thermal component, with temperatures of 5--10 MK and equivalent radii of 1--5 km. The current status of AXPs, including their comparison with SGRs, is reviewed by S.\\ Mereghetti et al.\\ (this volume). From the observational perspective, the nature of these enigmatic objects can hardly be understood without deep multiwavelength observations, particularly in the IR and hard X-ray bands, which would help elucidate the origin of the nonthermal component and separate the thermal component more precisely. We did not touch ``elderly'' nearby pulsars (e.g., B1929--10, B0950+08), mostly because they have not been observed with \\cha. If future deep observations show that their X-ray radiation is indeed thermal (emitted from hot, small polar caps), as suggested by the \\ros\\ and \\asc\\ observations, the analysis of the spectra and light curves will help establish the strength and geometry of their magnetic fields and the NS mass-to-radius ratio. We did not mention important non-detections of thermal radiation from very young pulsars, the Crab (Tennant et al.\\ 2001; M.C.\\ Weisskopf, this volume) and the recently discovered PSR J0205+6449 in the SNR 3C\\, 58 (Slane et al.\\ 2002), which put upper limits on the NS surface temperature. Finally, very interesting results on thermal radiation from transiently accreting NSs in quiescence have been reported recently (Rutledge et al.\\ 2001, and references therein). Although these old NSs are not truly isolated, their radiation in quiescence resembles that of young isolated NSs because of additional heating due to pycnonuclear reactions in the compressed accreted material.  The \\cha\\ observations described above have shown a few important things. First of all, they left no doubts that thermal emission indeed dominates in the soft X-ray radiation of quite a few isolated NSs, and the analysis of this radiation allows one to measure the NS temperatures\\footnote{See Fig.\\ 5 in Mereghetti et al.\\ (this volume) which demonstrates current estimates for the temperatures and luminosities of 22 thermally emitting NSs.} and, in some cases, the radii or mass-to-radius ratios. Particularly important is the conclusion that the simple picture of a NS with a centered dipole magnetic field and uniform surface temperature is, most likely, an oversimplification. We see that whenever the data allow a detailed analysis, the temperature is not uniform, in both active pulsars and radio-quiet NSs. Moreover, the radius of X-ray emitting area is, as a rule, considerably smaller than the ``canonical'' NS radius. The example of J1210--5226 (Sec.\\ 3.1) shows that the characteristic age ($\\tau_c=P/2\\dot{P}$) of a pulsar can differ from its true age by a large factor, and the conventional ``pulsar magnetic field'' can be quite different from the actual magnetic field at the NS surface. Therefore, any inferences on the properties of NSs and the superdense matter obtained without taking this into account should be considered with caution. For example, adopting the X-ray emitting radius to be the NS radius, a number of authors hurried to conclude that NSs are strange quark stars. Another example is the numerous comparisons of the NS cooling curves with observations of thermal X-rays from pulsars postulating that the pulsar's characteristic age is the true age and the empirical temperature of one of the ``thermal components'' is a single effective temperature. If we want to understand the real objects, we should abandon the naive paradigm and answer the basic questions raised by the observations: \\linebreak $\\bullet$ What is the true explanation of the small, hot areas in many, if not all, NSs, such as J2323+5848 (CCO of Cas A), RX J1856.5--3754, etc? If these small heated areas are associated not with (unseen) pulsar activity, but with superstrong local magnetic fields, how such fields affect the thermal evolution of NSs? \\linebreak $\\bullet$ What is the relation between the various kinds of NSs --- CCOs, AXPs, SGRs, active pulsars, dim NSs? For instance, why do CCOs, AXPs and SGRs show similar properties of thermal radiation, being so different in the other observational manifestations? Can it be that the members of different subclasses are intrinsically similar objects at different stages of their evolution (e.g., AXPs are descendants of CCOs) or viewed from different directions (e.g., dim NSs are usual middle-aged pulsars whose pulsar radiation is not seen because of unfavorable geometry)? \\linebreak $\\bullet$ What is the actual temperature distribution over the NS surface? Does the two-component model for thermal radiation we have used (thermal soft + thermal hard) describe the distribution adequately? \\linebreak $\\bullet$ What is the actual mechanism(s) of the surface emission? Are there any gaseous atmospheres at all? Why does the blackbody model describes the observed spectra so well in many cases?\\\\ These are the questions to be addressed (and, hopefully, answered) by future observations of NSs."
    },
    "0206/astro-ph0206354_arXiv.txt": {
        "abstract": "We investigate the accuracy and reliability of the semiempirical period - blue amplitude - V-band luminosity relationship for ab-type RR Lyrae stars originally obtained by Castellani \\& De Santis (1994) and De Santis (1996). We infer that the zero point of this relationship does depend on the metallicity, by studying a sample of both field and cluster variables. We also show that the use of this relationship can still be useful for those stellar systems showing an intrinsic metallicity spread, since in this case the metallicity effect has a negligible effect on the final distance modulus estimate.  We compare the adopted semiempirical relationship with the fully empirical one recently provided by Kov\\'acs \\& Walker (2001). When the zero point of the latter relation is fixed consistently with the former one, the two equations are equivalent.  By appling the semiempirical period - blue amplitude - V-band luminosity relation, as well as the technique proposed by Cassisi, De Santis \\& Piersimoni (2001), to the globular cluster $\\omega$~Cen, we show that the empirical slope of the relationship between the mass of the fundamental RR Lyrae pulsators and their metallicity, is in fair agreement with the one predicted by updated evolutionary models for Horizontal-Branch stars. ",
        "introduction": "The traditional distance indicator for Population II stellar systems is the magnitude of RR Lyrae variables; for this reason, several observational and theoretical investigations have been devoted to this class of variables (Bono, Castellani \\& Marconi 2000, Clement et al. 2001).  In spite of the large body of work devoted to their study, some relevant questions are still unanswered. One of the most important problems is the lack of a general agreement about both the slope and the zero-point of the absolute magnitude - metallicity ($M_V(RR) - [Fe/H]$) relationship, characteristic of the RR Lyrae stars (e.g., Caputo 1997, Gratton 1998, Cassisi, De Santis\\footnote{We inform with great sorrow that R. De Santis died on the 4th March 2002} \\& Piersimoni 2001, Benedict et al. 2002 and references therein). Some methods, like Baade-Wesselink and statistical parallax analyses applied to field RR Lyrae stars, plus observations of field Horizontal Branch (HB) stars with parallax measurements (but with large parallax errors --  Gratton 1998) support the {\\sl short} distance scale. On the other hand, the pulsational properties of cluster RR Lyrae stars (Sandage 1993), the main-sequence fitting to local subdwarfs (Gratton et al. 1997), the calibration of HB luminosity obtained by using the Cepheid distance modulus of the Large Magellanic Cloud (Walker 1992), and analysis based on double-mode RR Lyrae (Kov\\'acs 2000, Popielski, Dziembowski \\& Cassisi 2000) support the {\\sl long} scale.  One possible explanation of such a disagreement could be the existence of a true luminosity difference between field and cluster HB stars, as suggested by Gratton (1998) on the basis of an Hipparcos calibration of the absolute magnitude of field HB stars. However, this evidence is not supported by the analyses performed by Catelan (1998), De Santis \\& Cassisi (1999, hereinafter DC) and Carretta, Gratton \\& Clementini (2000). The origin of the distance dichotomy is therefore still unexplained and it has a big impact on a wide range of astrophysical problems such as globular clusters (GCs) age determinations and the extragalactic distances measurements.  A significant contribution to the solution of this problem can be provided by the analysis of the pulsational properties of RR Lyrae stars both in the Galaxy and in the Large Magellanic Clouds. A first step towards this direction was made by Sandage, Katem \\& Sandage (1981) who suggested the existence of a tight correlation between temperature and amplitude of RR Lyrae stars. More recently, Caputo \\& De Santis (1992) showed the existence of a clear correlation between period, blue amplitude and light-mass ratio of the variables. The reliability of these and similar relationships is important, since period and amplitude can be measured with high accuracy, regardless of uncertainties on both distance modulus and reddening. In this field, a pivotal importance is played by any relation connecting pulsational properties to the intrinsic luminosity of RR Lyrae stars.  Theoretical support for the existence of a correlation between visual magnitude, period and blue amplitude (hereinafter {\\it PLA} relation) for fundamental RR Lyrae ($RR_{ab}$) pulsators has been presented by Castellani \\& De Santis (1994) and by De Santis (1996). They provided also a semiempirical calibration of this relationship by adopting theoretical pulsational models and the observational database available at that time. However, until now, we lack a definitive assessment of the reliability of such relationship.  This paper is the third of a series investigating how the pulsational properties of RR Lyrae stars can be used to constrain their intrinsic luminosity. In particular, DC have used the pulsational behaviour of $RR_{ab}$ stars to obtain an accurate estimate of the absolute bolometric luminosity of ZAHB stars in GCs. It is worth noticing that their results do not depend on the underlying evolutionary models of HB stars, and this occurrence allowed them to perform a significant comparison with recent theoretical evaluations of the ZAHB luminosity. This, in turn, is important in order to properly evaluate the reliability of current theoretical models of low-mass, He-burning stars (see also the discussion in Vandenberg et al. 2000) .  Cassisi, De Santis \\& Piersimoni (2001, hereinafter CDP) have adopted the same method outlined by DC, in order to derive the absolute visual magnitude of the ZAHB within the RR Lyrae instability strip, for a sample of galactic GCs with accurate photometric data for both variable and non-variable HB stars. After applying a correction for the difference between the mean RR Lyrae magnitude and the ZAHB one, they derived a  $<M_V(RR)> - [Fe/H]$ relation and compared it with the most recent empirical ones (see also Caputo et al. 1999). They also discussed a method for determining the GCs distance based only on the pulsational properties of their RR Lyrae population. The advantage of this method is that it does not need an estimate of the ZAHB level, which is particularly difficult task in the case of GCs with blue HB, whose RR Lyrae are suspected to be evolved stars.  The approaches developed by DC and CDP require a preliminary estimate of the cluster metallicity in order to determine the appropriate mass range for $RR_{ab}$ stars. Due to the non-negligible uncertainties affecting both the GC metallicity scale (see, i.e., Rutledge et al. 1997 and Vandenberg et al. 2000) and their $\\alpha$-elements distribution, this reduces the accuracy of the method. Therefore, in the present work we investigate the possibility to estimate the absolute visual magnitude of $RR_{ab}$ stars by adopting a magnitude-period-amplitude relation. This would not require a preliminary evaluation of the stellar metallicity.  In the next section we briefly review the semiempirical PLA relationship adopted in present analysis. In section 3, we apply this relationship to a sample of field RR Lyrae stars, in order to investigate on its accuracy when applied to single stars. In section 4, we follow the same approach in case of a selected sample of galactic GCs, and the derived distance moduli are compared with those provided by CDP. In section 5, the semiempirical PLA relation is compared with the empirical one by Kovacs \\& Walker (2001). An application of the method to stellar systems showing a spread in the metallicity, like the GC $\\omega$~Cen, is shown in section 6. A brief discussion and conclusions follow in the last section. ",
        "conclusions": "By applying both the PLA relation and the CDP method to the GC $\\omega$~Cen, we have shown that the two approaches are in agreement when applied to stellar systems with an internal metallicity spread. This is an advantage of the PLA relationship over the CDP method, since the former does not require any preliminary metallicity estimate. In this respect, some dwarf galaxies in the Local Group (LG) are a fine target to exploit this technique. In fact, even if current metallicity measurements, based mainly on photometric indices, reveal an intrinsic metallicity dispersion of the order of 0.3-0.5 dex (Mateo 1998) in most of the LG dwarfs, more recent high-resolution spectroscopic measurements have disclosed the existence of larger internal spread in the heavy elements abundance, for instance, $\\Delta[Fe/H]=0.73$ dex and 1.53 have been obtained by Shetrone, Cote \\& Sargent (2001) for Ursa Minor and Draco, respectively. \\begin{figure} \\psfig{figure=fig9p.ps,height=13cm,width=10cm} \\caption{Comparison between the apparent visual magnitude and the $M_V$ value obtained from the PLA relation, for the RR Lyrae variables in the dwarf galaxies Leo II, Sculptor and Sextans.} \\end{figure}  We have applied the PLA relation to the RR Lyrae sample in three LG dwarfs: Leo II, Sculptor and Sextans. The data are from Siegel \\& Majewski (2000) for Leo II, from Kaluzny et al (1995) for Sculptor, and from Mateo, Fisher \\& Krzeminski (1995) for Sextans. For each dwarf we have measured the distance modulus by applying the PLA relation to each individual variable and then computed the mean of the various values.  The resulting distances are $21.67\\pm0.05$ mag, $19.60\\pm0.03$ mag, and $19.83\\pm0.06$ mag for Leo II, Sculptor and Sextans, respectively (the listed error corresponds to the maximum random error, so it does not account for the observational uncertainty affecting the apparent visual magnitudes). These estimates are in good agreement with currently adopted values (Mateo 1998). In figure 9, we show the satisfactory level of agreement between the apparent visual magnitude for each RR Lyrae variable and the $M_V$ value provided by the PLA relation.  Since the systematic metallicity effect, affecting the PLA relation, works in opposite directions, there is a good agreement between the distance moduli derived by using the CDP or PLA methods. In the case of a single metallicity stellar population, a reliable distance measurement should be still achieved. In fact, the maximum systematic error affecting the GC distance moduli provided by the PLA relation is of the order of $\\pm0.04$ mag.  However, the accuracy of the absolute magnitudes provided by the PLA relation strongly relies on the reliability of the adopted blue amplitude for the $RR_{ab}$ variables. This means that one has to exclude from the adopted data sample those variables affected by amplitude modulation ({\\sl Blazhko} effect).  We wish to summarize the assumptions upon which the PLA relation zero point relies:  \\begin{itemize} \\item{the validity of the adopted temperature scale. Concerning this topic, De Santis (2001) has shown that this relation does not significantly depend on the selected sample of variables, and has discussed its range of validity. Moreover, DC have also shown that, when this pulsational temperature scale is used, consistency is achieved between several sets of $T_{eff}$-color transformations (Buser \\& Kurucz 1978, Bessel, Castelli \\& Plez 1998, and Green 1988) and the Lub's reddening scale. However, a significant error in the zero-point of this reddening scale and/or of the above quoted transformations, would translate into an error for the zero-point of equation (3);}  \\item{in section 1, we have noted that the use of the van Albada \\& Baker (1971) relation for the fundamental pulsational equation makes the PLA relation fainter by 0.05 mag in comparison with the Bono et al's (1997) relation;}  \\item{when calibrating the zero point of the PLA relation, DS adopted $0.69M_\\odot$ as the mean mass of fundamental pulsator for a metallicity $[Fe/H]=-1.5$, while this assumption seems to be supported by recent HB stellar models, it relies on the reliability of the stellar model input physics.} \\end{itemize}  \\noindent Since all the quoted assumptions seem to be reliable, we trust the accuracy of the semi-empirical PLA relation for deriving the distance to stellar systems such as galactic GCs and LG dwarf galaxies.  {\\bf Acknowledgments:}  We warmly thank F. Caputo and M. Salaris for an accurate reading of this manuscript and for the suggested improvements. We also wish to thank G. Kov\\'acs and A.R. Walker for their useful comments and suggestions on an early draft of this paper. Financial support for this work was provided by MURST-Cofin 2002."
    },
    "0206/astro-ph0206448_arXiv.txt": {
        "abstract": "Inspired by recent developments in particle physics, the so-called brane world cosmology seems to provide an alternative explanation for the present dark energy problem. In this paper, we use the estimated age of high-$z$ objects to constrain the value of the cosmological parameters in some particular scenarios based on this large scale modification of gravity. We show that such models are compatible with these observations for values of the crossover distance between the 4 and 5 dimensions of the order of $r_c \\leq 1.67H_o^{-1}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206162_arXiv.txt": {
        "abstract": "We present a rigorous, detailed study of the generic, quantitative properties of gravitational lensing near cusp catastrophes. Concentrating on the case when the individual images are unresolved, we derive explicit formulas for the total magnification and centroid of the images created for sources outside, on, and inside the cusped caustic.  We obtain new results on how the image magnifications scale with respect to separation from the cusped caustic for arbitrary source positions.  Along the axis of symmetry of the cusp, the total magnification $\\mu$ scales as $\\mu \\propto u^{-1}$, where $u$ is the distance of the source from the cusp, whereas perpendicular to this axis, $\\mu \\propto u^{-2/3}$.  When the source passes through a point $\\bu_0$ on a fold arc abutting the cusp, the image centroid has a jump discontinuity; we present a formula for the size of the jump in terms of the local derivatives of the lens potential and show that the magnitude of the jump scales as $|u^0_1|^{1/2}$ for $|u^0_1| \\ll 1$, where $|u^0_1|$ is the horizontal distance between $\\bu_0$ and the cusp.  The total magnifications for a small extended source located both on, and perpendicular to, the axis of symmetry are also derived, for both uniform and limb darkened surface brightness profiles.  We find that the difference in magnification between a finite and point source is $\\la 5\\%$ for separations of $\\ga 2.5$ source radii from the cusp point, while the effect of limb-darkening is $\\la 1\\%$ in the same range. Our predictions for the astrometric and photometric behavior of both pointlike and finite sources passing near a cusp are illustrated and verified using numerical simulations of the cusp-crossing Galactic binary-lens event MACHO-1997-BUL-28.  Our results can be applied to any microlensing system with cusp caustics, including Galactic binary lenses and quasar microlensing; we discuss several possible applications of our results to these topics. ",
        "introduction": "Over the past twenty years, gravitational lensing has grown from a mere curiosity to an important component of a large and diverse set of fields in astronomy.  Its ubiquity is due at least in part to the fact that its effects are observable over a wide range of scales. This has enabled astronomers to use lensing to study everything from the smallest compact objects, to the largest structures in the universe, and almost everything in between. Despite the diversity of applications of gravitational lensing, the mathematical description of the phenomenon itself is both relatively tractable, and universal.  In almost all cases, gravitational lensing can be described by a two-dimensional mapping from a lens plane to light source plane.  Once this mapping is specified, all of the properties of a gravitational lens can be derived in principle.  The {\\it observable} properties of lensing, however, depend on the phenomenon to which it is applied.  Therefore, lensing is traditionally divided into an number of different regimes, which are delimited by the observables.  For example, the term microlensing is typically applied to the case when multiple images occur, but are not resolved.  When multiple images are created by a gravitational lens, the separation between these images is typically of order the Einstein ring radius, \\begin{equation} \\thetae=\\sqrt{{{4 G M}\\over c^2 D} }, \\label{eqn:thetae} \\end{equation} where $M$ is the mass of the lens, $D\\equiv \\dos\\dol/\\dls$, and $\\dos$, $\\dol$, and $\\dls$ are the distances from observer to source, observer to lens, and lens to source, respectively.  Thus, the term microlensing is applied when $\\thetae$ is less than the resolution.  In this case, all one can measure is the collective behavior of all the images created by the lens, i.e.,\\ the total magnification, and the position of the center-of-light (centroid) of the image.  In fact, if the observer, source, and lens were not in relative motion, than the individual image magnifications and positions would be fixed, and even these properties would not be measurable.  The relative positions of the observer, lens, and source, and thus the magnification and centroid, are expected to vary on time scales 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.  Fortunately, in the two regimes where microlensing has been discussed, the typical values of $\\thetae$, $\\dol$, and $v_\\perp$ result in reasonable time scales.  Typical values for the lens mass $M$, relative source-lens distance $D$, transverse velocity $v_\\perp$, and the resulting typical values for $\\thetae$ and $\\te$ for both Local Group and cosmological microlensing are given in Table 1.  Also shown are typical values for the radius  $\\theta_*$ of the source's emission region, this radius in units of $\\thetae$, $\\rho_*\\equiv \\theta_*/\\thetae$, and the time it takes the lens to cross the source, $t_*\\equiv \\rho_* \\te$. These latter parameters will be relevant to the discussion of finite source effects in \\S\\ref{sec:fsources}.  For the Local Group, a typical Einstein radius crossing time is $\\te={\\cal O}(100~{\\rm days})$, whereas for cosmological microlensing, $\\te={\\cal O}(10~{\\rm years})$.  \\begin{table}[t] \\begin{center} \\begin{tabular}{c|cccccccc} \\tableline \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$M$} & \\multicolumn{1}{c}{$D$} & \\multicolumn{1}{c}{$\\thetae$} & \\multicolumn{1}{c}{$v_\\perp$} & \\multicolumn{1}{c}{$\\te$} & \\multicolumn{1}{c}{$\\theta_*$} & \\multicolumn{1}{c}{$\\rho_*$} & \\multicolumn{1}{c}{$t_*$} \\\\ \\tableline Local Group  & $1 M_\\odot$ & $10~{\\rm kpc}$ & $1{\\rm mas}$ & $100~{\\rm km~s^{-1}}$ & $100~{\\rm days}$ & $1\\muas$ & $10^{-3}$ & $0.1~{\\rm days}$\\\\ \\tableline Cosmological  & $1 M_\\odot$ & $1~{\\rm Gpc}$ & $3\\muas$ & $500~{\\rm km~s^{-1}}$ & $10~{\\rm years}$ & $0.1\\muas$ & $0.03$ & $100~{\\rm days}$\\\\ \\tableline \\end{tabular} \\end{center} \\tablenum{1} {\\bf Table 1} Typical Microlensing Parameters. \\\\ \\label{tbl:table1} \\end{table}  Of exceptional importance in microlensing is the existence of caustics:  positions in the light source plane corresponding to the critical values of the lens mapping. On caustics, at least one image is formed that has formally infinite magnification (for a point-source).  When a source crosses a caustic, both the total magnification and centroid of all the images exhibit instantaneous, discontinuous jumps.  These jumps are averaged out over the finite source size; however, it is generically true that large gradients in the magnification and centroid exist near caustics.  Furthermore, microlensing caustics have several important and useful properties.  First, the large magnification results in a large photon flux from the source.  Second, the large gradient in the magnification and centroid with respect to source position effectively implies high angular resolution.  Finally, the highly localized nature of the high-magnification and large centroid-shift regions created by caustics results in characteristic, and easily-recognizable features, in both astrometric and photometric microlensing curves.  Many authors have suggested exploiting these properties of caustics to study a number of astrophysical applications, i.e.,\\ stellar multiplicity \\citep{mandp1991}, stellar atmospheres \\citep{gould2001}, individual microlens mass measurements \\citep{gandg2002}, microlens mass functions \\citep{wwt2000a}, properties of the emission regions of quasars \\citep{wps1990,ak1999,fw1999,wwt2000b}, and lens transverse velocities \\citep{wwt1999}.  See Gaudi \\& Petters (2002; hereafter Paper I) for a more thorough discussion of the uses of astrometric and photometric microlensing observations in the presence of caustics in both Local Group and cosmological contexts.  Although the caustic curves of gravitational lenses exhibit an enormously rich and diverse range of properties, it can be shown rigorously that each stable lensing map has only two types of caustic singularities: folds and cusps (Petters, Levine, \\& Wambsganss 2001, p. 294).  Each of these two types of singularities have generic and universal properties and, in particular, each can be described by a polynomial mapping from the lens plane to light source plane. The coefficients of these mappings depend on local derivatives of the dimensionless surface potential of the lens.  In Paper I, we used the mapping for a fold singularity to derive the observable properties of gravitational lensing near folds, paying particular attention to the case of microlensing, in which the images are unresolved. We derived analytic expressions for the total magnification and centroid shift near a generic, parabolic, fold caustic.  We then showed how these expressions reduce to those for the more familiar linear fold, which lenses a nearby source into two equal magnification, opposite parity images whose total magnification is proportional to $u^{-1/2}$, where $u$ is the distance of the source to the fold caustic.  We then generalized these results to finite source sizes.  Finally, we compared our analytic results to numerical simulations of the Galactic binary-lens event OGLE-1999-BUL-23, in which the source was observed to cross a fold caustic.  We found excellent qualitative agreement between our analytic and semi-analytic expressions for the photometric and astrometric behavior near a fold caustic, and our detailed numerical simulations of the second fold caustic crossing of OGLE-1998-BUL-23.  In this paper, we present a similarly detailed study of the generic, quantitative properties of microlensing near cusps.  Although fold caustic crossings are expected and observed to dominate the sample of caustic crossings in Galactic binary events \\citep{gandg1999,ggh2002, alcock2000}, cusp crossings will nevertheless represent a non-negligible fraction of all caustic crossing events.  In fact, at least two cusp crossing events have already been observed, the Galactic bulge events \\event ~\\citep{albrow1999a}, and MACHO-1997-BLG-41 \\citep{alcock2000, albrow2000}.  It is interesting to note that the analyses and modeling of these events were performed using entirely numerical methods.  As we discuss in some detail (see \\S\\ref{sec:apply}), we believe that the analytic results derived here are particularly amenable to the analysis of \\event~ and similar events.  In the cosmological context, the role of cusps versus folds is less clear, due primarily to the more complicated structure of the caustics themselves.  However, it has been shown that, in the limit of high magnifications, the inclusion of cusps alters the form of the total microimage magnification cross section, due to the lobe of high-magnification close to and outside the cusped caustic \\citep{sw92}.  Lensing near cusps has been substantially less well-studied than lensing near folds, and as a result, useful analytic expressions for the observable properties are few.  Previous studies have focused almost exclusively on the magnification of the images created by a cusp singularity \\citep{sw92, mao92, zak95, zak99}.  Here we study all the observable properties of gravitational microlensing near cusps, including the photometric (total magnification) and astrometric (light centroid) behavior, for both point sources and extended sources with arbitrary surface brightness profiles.   The layout of this paper is as follows.  In \\S\\ref{sec:analytic}, we derive analytic expressions for the image positions, magnification, and light centroid for sources near a generic cusp.  In \\S\\ref{sec:general} we define the observable microlensing properties. In \\S\\ref{sec:local}, we start with the generic expression for the mapping near a cusp, and derive all the properties for the local images for sources exterior to (\\S\\ref{sec:exterior}), on (\\S\\ref{sec:on}), and interior to (\\S\\ref{sec:interior}) the caustic. We generalize the discussion to include images not associated with the cusp in \\S\\ref{sec:global}, and study extended sources in \\S\\ref{sec:fsources}.  In \\S\\ref{sec:mb9728}, we illustrate the observable behavior near a cusp by numerically simulating the Galactic binary-lens cusp-crossing event \\event, and directly compare these numerical results with our analytic expressions in \\S\\ref{sec:compare}.  In \\S\\ref{sec:apply}, we suggest several possible applications in both local group and cosmological contexts.  We summarize and conclude in \\S\\ref{sec:summary}. We note that, for the sake of completeness, we include some results that have been presented elsewhere.  Combined with the results from Paper I, the results presented here describe the observable properties of gravitational microlensing near all stable singularities. ",
        "conclusions": "}  We have presented a comprehensive, detailed, and quantitative study of gravitational lensing near cusp catastrophes, concentrating on the specific regime of microlensing (when the individual images are unresolved).  We started from a generic polynomial form for the lens mapping near a cusp that relates the image positions to the source position. This mapping is valid to third order in the image position.  The quantitative properties of this mapping are determined by the polynomial coefficients, which can be related to local derivatives of the projected potential of the lens.  Near a cusp, the critical curve is a parabola, which maps to the cusped caustic.  We find an simple expression for the vertical component of the image position $\\theta_2$, which is a cubic of the form $\\theta_2^3 + \\ttp \\theta_2 + \\ttq=0$, where $\\ttp$ and $\\ttq$ are functions of the source position $\\bu$.  The solutions of the cubic in $\\theta_2$ are characterized by the discriminant $D(\\bu)=(\\ttp/3)^3 + (\\ttp/2)^2$.  For source positions $\\bu$ outside the caustic, we have $D(\\bu)>0$ and, thus, there is locally one image.  We determined the magnification and location of this image, and showed, in particular, that along the axis of the cusp (tangent line to the cusp), the magnification scales as $u^{-1}$, where $u$ is the distance from the cusp point. Perpendicular to the axis, the magnification scales as $u^{-2/3}$.  We also determined the image positions and magnifications for sources on the caustic, where $D(\\bu)=0$.  On the caustic, there are two images.  One image is infinitely magnified and results from the merger/creation of a pair of images, whereas the second image has finite magnification, and can be smoothly joined to the single lensed image of a source just outside the caustic.  For sources inside the caustic ($D(\\bu)<0$), we find that there are three images.  One image has positive parity and diverges as the source approaches the top fold caustic abutting the cusp, but can be continuously joined to the non-divergent image when the source is on the bottom fold.  Similarly, the other positive parity image diverges as the source approaches the bottom fold, but can be continuously joined to the non-divergent image when the source is on the top fold.  The third image has negative parity, and diverges as the source approaches either fold caustic.  All three images diverge as one approaches the cusp point.  For sources on the axis of the cusp, but interior to the caustic, the total magnification of all the images diverges as $u^{-1}$.  We also derived analytic expressions for the centroid of all three images created when the source is interior to the caustic.  We generalized our results beyond the local behavior near the cusp by deriving expressions for the total magnification and centroid including the images not associated with the cusp.  We further considered rectilinear source trajectories, and parameterized this trajectory in order to calculated the dependence of the photometric and astrometric behavior on time.  In particular, due to the presence of the infinitely magnified images when the source is on the fold caustics, but finite magnification outside the caustic, we find that the centroid exhibits a finite, instantaneous jump whenever the source crosses one of the two folds abutting the cusp.  We present a formula for the magnitude of the jump that depends only on the local coefficients of the cusp mapping, and the location of the caustic crossing.  We note that this magnitude decreases monotonically as a function of the horizontal distance between the cusp and where the source crosses the cusped caustic curve, so that a source which crosses the cusp point exactly exhibits no jump discontinuity.   Beginning with the appropriate modifications to the formulae for the total magnification and centroid for finite sources with arbitrary surface brightness profiles, and combining these with the analytic results we obtained for the magnification and centroid for point sources near a cusp, we derived semi-analytic expressions for the uniform and limb-darkened finite source magnification for small sources on and perpendicular to the axis of the cusp.  We also derived expressions for the centroid of a small source on the axis for uniform and limb-darkened sources.  In order to illustrate the photometric and astrometric lensing behavior near cusps, and to provide order-of-magnitude estimates for the effect of finite sources and limb-darkening on these properties, we numerically calculated the total magnification and centroid shift for the observed cusp-crossing Galactic binary microlensing event \\event.  We find that the cusp crossing results in large, ${\\cal O}(\\thetae)$ centroid shifts, which should be easily detectable with upcoming interferometers.  We find that limb-darkening induces a deviation in the centroid of $\\sim 35\\Gamma \\muas$.  We compared our numerical calculations with the analytic expectation, and found excellent agreement.  Adjusting only the magnification and centroid of the images unrelated to the cusp, and adopting the coefficients appropriate to the cusp of \\event, we find that our analytic formulae predict the magnification to $\\la 5\\%$ and the centroid to $5\\muas$, for positions within $\\sim 2$ source radii of the cusp.  Finally, we suggested several applications of our results to both Galactic and cosmological microlensing applications.  We suggest that one can use our analytic expressions to determine the applicability of the point-source approximation for sources near a cusp.  In particular, we note that the finite source magnification deviates from the point-source magnification by $\\la 5\\%$ for sources separated by $\\ga 2.5$ source radii.  We also discussed how the local and generic behavior of the cusp can be used to simplify the fitting procedure for cusp-crossing events.  Lastly, we outlined a method by which the typical angular Einstein ring radius of the perturbing microlenses of a macrolensed quasar might be estimated using measurements of the jump in the centroid that occurs when the source crosses a fold, making use of the analytic expression for the magnitude of the jump derived here.  Despite their apparent diversity, the mathematical underpinning of all gravitational lenses is identical.  In particular, all lenses exhibit only two types of stable singularities: folds and cusps.  In Paper I, we studied gravitational microlensing near fold caustics; here we have focussed on microlensing near cusp caustics.  A generic form of the mapping from source to image plane near these each of these types of caustics can be found, and used to derive mostly analytic expressions for the photometric and astrometric behavior near folds and cusps. These expressions can be used to predict the behavior near all stable caustics of gravitational lenses, and applied to a diverse set of microlensing phenomena, including Galactic binary lenses and cosmological microlensing."
    },
    "0206/astro-ph0206481_arXiv.txt": {
        "abstract": "\\noindent  We present our investigation on the effect of warps on the extraction of rotation curves in edge-on galaxies. The method to derive the rotation curve from the position-velocity diagram in warped edge-on systems yields underestimated velocities, and the tilted-ring model is not reliable in highly inclined, poorly resolved galaxies. In a warped system the kinematical major axis is different from the optical major axis. While this is generally a limit in optical slit spectroscopy, in the \\HI\\ emission which extends far from the optical body where self-gravity is weaker and the effect of warping is more pronounced, this represents a severe effect to be considered in the procedure to extract the rotation curve.  We propose a new approach to extract the rotation curve in highly inclined, warped galaxies. Based on this method we are able to trace accurately the frequency of peculiarities in our sample of Thick Boxy Bulge (TBB) galaxies. We report an increasing trend of kinematical lopsidedness from spheroidal bulge galaxies towards TBB galaxies.  Concerning the question whether interactions contribute significantly to the bar formation and to the subsequent evolution in a box/peanut (b/p) structure, we confirm these theoretical predictions. Based on our sample, galaxy interaction is the likely formation mechanism to trigger bars in TBB galaxies. ",
        "introduction": "A significant fraction of bulges -- up to 50\\% -- deviate from the classical spheroidal shape, and display excess light above and below the galactic disks, resembling a boxy or peanut shape when viewed edge-on. For this class of bulges there is general agreement that the box/peanut (b/p) shape of the light distribution is related to a bar. This scenario is strongly supported by observations of the stellar kinematics (Bureau \\& Freeman 1999) as well as by more recent N-body simulations (Athanassoula et al. 2002). However, the origin of such bars is still not well understood (Sellwood 2000).  At present, two scenarios are widely assumed for the bar formation. The first one predicts bar formation through the spontaneous instabilities in relatively cool and rotationally supported disks, while in the other one the bar is triggered by interactions. In fact, interactions both in the form of tidal interaction and direct interaction (minor merger) can provide a very efficient mechanism to trigger instabilities in the disks, and initiate or speed up bar formation. Once a bar is formed, it is free to evolve in a b/p structure due to buckling and thickening (Noguchi 1987; Gerin, Combes \\& Athanassoula 1990; Mihos et al. 1995).  The objects of our investigation are b/p bulge galaxies, which however show remarkable differences compared to the classical b/p-shaped bulge galaxies in the prominence and thickness of their b/p structures. Because of these morphological features, we will refer to them as Thick Boxy Bulges galaxies (TBBs, L\\\"utticke et al. 2000).  If interactions, spontaneous instabilities, or intermediate mechanisms are the reasons for the origin and the evolution of these TBB galaxies they may leave different traces which should be observed.  For this purpose we have selected a sample of 8 TBB galaxies, and we are investigating the likely formation mechanism of this class of objects. The galaxies  were selected from the Third Reference Catalogue of Bright Galaxies (de Vaucouleurs et al. 1991) on the basis of their orientation (nearly edge-on) and their diameters. The diameters of the galaxies are constrained to be larger than 2$'$ at the B$_{25}$ isophote on the Digital Sky Survey (DSS).  Based on the fact that interaction events produce a highly asymmetric distribution and complex kinematics of the gaseous component, large deviations from symmetry in the morphology and kinematics will indicate a galaxy which is strongly interacting with nearby companions. In order to discern different bar formation scenarios (spontaneous instabilities or bar triggered by interactions) one of the diagnostics is the frequency of peculiarities of the TBB galaxies compared to spheroidal galaxies.  We are studying the statistics of warps, their shapes as well as lopsidedness. The latter is investigated both in the kinematics (comparing the approaching and receding side of the rotation curve) and in the density distribution (comparing the mass included in the receding/approaching side) using 21-cm line VLA and ATCA data. ",
        "conclusions": ""
    },
    "0206/astro-ph0206432_arXiv.txt": {
        "abstract": "The SCUBA 8-mJy survey is the largest submillimetre (submm) extragalactic mapping survey undertaken to date, covering 260\\,arcmin$^2$ to a 4\\,$\\sigma$ detection limit of $\\simeq$\\,8\\,mJy at 850\\,$\\mu$m, centred on the Lockman Hole and ELAIS N2 regions. Here, we present the results of new 1.4-GHz imaging of these fields, of the depth and resolution necessary to reliably identify radio counterparts for 18 of 30 submm sources, with possible detections of a further 25 per cent. Armed with this greatly improved positional information, we present and analyse new optical, near-infrared (IR) and {\\it XMM-Newton} X-ray imaging to identify optical/IR host galaxies to half of the submm-selected sources in those fields. As many as 15 per cent of the submm sources detected at 1.4\\,GHz are resolved by the 1.4$''$ beam and a further 25 per cent have more than one radio counterpart, suggesting that radio and submm emission arise from extended starbursts and that interactions are common. We note that less than a quarter of the submm-selected sample would have been recovered by targeting optically faint radio sources, underlining the selective nature of such surveys. At least 60 per cent of the radio-confirmed optical/IR host galaxies appear to be morphologically distorted; many are composite systems --- red galaxies with relatively blue companions; just over one half are found to be very red ($I-K>\\rm 3.3$) or extremely red ($I-K>\\rm 4$); contrary to popular belief, most are sufficiently bright to be tackled with spectrographs on 8-m telescopes.  We find one submm source which is associated with the steep-spectrum lobe of a radio galaxy, at least two more with flatter radio spectra typical of radio-loud active galactic nuclei (AGN), one of them variable. The latter is amongst four sources ($\\equiv$\\,15 per cent of the full sample) with X-ray emission consistent with obscured AGN, though the AGN would need to be Compton thick to power the observed far-IR luminosity. We exploit our well-matched radio and submm data to estimate the median redshift of the $S_{\\rm 850\\mu m}$\\,$\\sim$\\,8\\,mJy submm galaxy population. If the radio/far-IR correlation holds at high redshift, and our sample is unbiased, we derive a conservative limit of $<\\!z\\!>$\\,$\\geq$\\,2.0, or $\\geq$\\,2.4 using spectral templates more representative of known submm galaxies. ",
        "introduction": "The nature of the sources detected in deep submm and mm surveys remains controversial. All SCUBA surveys agree as to the high surface density of 850-$\\mu$m sources detected at the mJy level (Smail, Ivison \\& Blain 1997; Hughes et al.\\ 1998; Barger, Cowie \\& Sanders 1999a; Eales et al.\\ 1999; Chapman et al.\\ 2002a; Borys et al.\\ 2002; Webb et al.\\ 2002b) but their exact distances, luminosities and their power source all remain contentious subjects.  Most of the far-IR/submm background detected by the {\\it DIRBE} and {\\it FIRAS} experiments (Puget et al.\\ 1996; Fixsen et al.\\ 1998; Hauser et al.\\ 1998; Schlegel, Finkbeiner \\& Davis 1998) has already been resolved into discrete sources by SCUBA (Blain et al.\\ 1999b; Smail et al.\\ 2002a; Cowie et al.\\ 2002) implying that the cosmic energy budget in the early Universe was dominated by hitherto undetected dust-enshrouded systems, either starbursts with star-formation rates $\\gg$100\\,M$_{\\odot}$\\,yr$^{-1}$, sufficient to construct a giant elliptical galaxy in $\\ls$1\\,Gyr, or Compton-thick AGN associated with the formation of super-massive black holes (SMBH).  If the submm galaxy population lies at high redshift, $z\\sim 3$, and is predominantly powered by star formation, then its star-formation rate density is higher than that deduced from optical/ultraviolet observations of the more numerous Lyman-break galaxies (Steidel et al.\\ 1999), a population with which there appears to be little overlap (Peacock et al.\\ 2000; Chapman et al.\\ 2000; Webb et al.\\ 2002a; cf.\\ Adelberger \\& Steidel 2000).  In this scenario, the properties of SCUBA galaxies (e.g.\\ space density, redshift distribution, etc.) would need to be reproduced by any successful model of galaxy formation.  Equally, if the bulk of the bolometric luminosity of this population derives from gravitational accretion onto black holes then they clearly represent a crucial phase in the formation of SMBH and the evolution of QSOs and powerful radio galaxies (Archibald et al. 2001; Page et al.\\ 2001).  The apparently tight relation seen locally between the masses of bulges and the those of their resident SMBH suggests that both of these scenarios may contain elements of truth, indicating a complex interplay between obscured star formation, AGN activity and feedback in the early evolution of spheroids and SMBH (Silk \\& Rees 1998; Fabian 1999; Archibald et al.\\ 2002).  While there has been significant progress in detailing the observational properties of the SCUBA population, theoretical interpretation has lagged behind. The standard framework for the theoretical understanding of this population relies upon hierarchical models which employ the cold dark matter (CDM) paradigm.  These have successfully described the properties of the galaxies and large-scale structure in the local Universe (e.g.\\ Cole et al.\\ 2000) but the gradual growth of the characteristic mass of galaxies leads these models to predict that the most massive galaxies have formed only recently, $z\\rm\\ls 1$ (Kauffmann \\& Charlot 1998), even in a low-density $\\lambda$CDM cosmology.  Semi-analytic models of galaxy formation, developed within the hierarchical framework, predict that these massive galaxies form primarily through mergers, where the attendant starburst activity can be sufficient to power the prodigious luminosities seen in local ultraluminous IR galaxies (ULIRGs --- Baugh et al.\\ 2001).  However, the strong decline in the number density of massive galaxies with redshift means that these models predict relatively modest median redshifts for the most massive mergers, $z\\rm\\ls 1$, unless the physical nature of the systems evolves radically (Blain et al.\\ 1999a, 1999c), or the efficiency of high-mass star formation is greater in bursts than in the quiescent mode seen in local disks. The most natural prediction of these models is therefore a low median redshift for galaxies selected by SCUBA. If it is shown that submm galaxies lie predominantly at high redshift, $z\\gg \\rm 1$, and that they represent massive gas-rich mergers (most probably associated with the formation epoch of massive ellipticals, Eales et al.\\ 1999) then this will require a radical overhaul of the treatment of high-redshift star formation in CDM-based hierarchical models. Hence an estimate of the redshift distribution, $N(z)$, for a complete, robust and well-characterised sample of submm-selected galaxies provides one important test of current theoretical galaxy formation models. In addition, the $N(z)$ is crucial for estimating the true 3-dimensional clustering of the submm population from the projected 2-dimensional clustering of sources in panoramic SCUBA surveys.  The strength of the clustering of submm galaxies reflects the mass (and bias) of these systems and provides a further test of the predictions from galaxy formation models.  For these reasons, determining the $N(z)$ of complete samples of submm galaxies is one of the highest priorities for researchers working on this enigmatic population (e.g.\\ Blain et al.\\ 1999c, 2000; Smail et al.\\ 2000, 2002a).  Unfortunately, if the majority of the submm population have no plausible optical counterparts, as has been widely reported, then traditional optical spectroscopy is not a viable option for determining $N(z)$ (e.g.\\ Barger et al.\\ 1999b). The faintness of near-IR counterparts to submm sources gives little hope to IR spectroscopists either and attention has focussed on redshift engines of one sort or another or on broadband photometric techniques (e.g.\\ Townsend et al.\\ 2001; Hughes et al.\\ 2002; Aretxaga et al.\\ 2002).  One potentially profitable route exploits the well-known radio/far-IR correlation (Dickey \\& Salpeter 1984; de Jong et al.\\ 1985; Helou, Soifer \\& Rowan-Robinson 1985) as a redshift estimator using deep radio observations of submm sources.  The submm flux density, $S_{\\nu}$, goes as $\\nu^{\\sim3.5}$, while for the optically thin synchrotron emission in the radio, $S_{\\nu} \\propto \\nu^{-0.7}$ (Condon 1992). $S_{\\rm 850\\mu m}/S_{\\rm 1.4\\,GHz}$ is thus a sensitive function of redshift, initially rising as $(1+z)^{\\gs 4}$ (Carilli \\& Yun 1999). Observations at 1.4\\,GHz thus complement submm surveys perfectly, being similarly sensitive to star-forming galaxies, although only at $z \\rm\\ls 3$ with present facilities (at $z\\gs 3$, the positive $K$ correction at 1.4\\,GHz overcomes the available sensitivity).  Given the preponderance of possible optical counterparts at the $I \\le 26$ level, the other crucial role of radio observations is to exploit their superior resolution to tie down the positions of submm sources: $\\sigma \\sim\\rm 0.3''$ compared to 4$''$ for SCUBA (e.g.\\ Ivison et al.\\ 1998, 2000b, 2001). Moreover, a single radio image can cover $\\sim$\\,500\\,arcmin$^2$ with high sensitivity and $\\sim$\\,1$''$ resolution (for 25-m antennas separated by $\\sim$\\,30\\,km at 1.4\\,GHz) enabling many of the sources in even the largest submm surveys to be identified in a single radio map. In addition, the large field of view allows the radio coordinate frame to be aligned accurately with the optical/IR frame (see \\S2.3). This means that only the positions of the most distant galaxies, those undetected in the radio, need then be laboriously determined on a case-by-case basis, via mm-wave continuum interferometry at the Owens Valley Radio Observatory (e.g.\\ Frayer et al.\\ 2000) and at Plateau de Bure (e.g.\\ Downes et al.\\ 1999; Lutz et al.\\ 2001).  Radio observations also act as a useful probe of AGN, regardless of the level of obscuration, via the identification of lobe-like morphologies or deviations of the radio spectral index ($\\alpha$, where $S_{\\nu} \\propto \\nu^{\\alpha}$) from the $-0.7$ expected for star-forming galaxies (e.g.\\ SMM\\,J02399$-$0136, Ivison et al.\\ 1999), or via anomalously high radio fluxes (e.g.\\ SMM\\,J14009+0252, Ivison et al.\\ 2000b).  Previous radio imaging of submm samples has been extremely successful, identifying robust optical/IR counterparts (Ivison et al.\\ 1998, 2000b, 2001; Smail et al.\\ 1999) and providing evidence that submm-selected galaxies are extremely distant, $z$\\,$\\ge$\\,2--2.5 (Smail et al.\\ 2000, 2002a; cf.\\ Lawrence 2001).  To date, however, the approach has been limited by small-number statistics, by the narrow, deep nature of the Smail et al.\\ (2002a) survey, which has a median lensing-corrected flux of $4.0\\pm 0.7$\\,mJy, and by the need to spread observing time across many fields (although this was mitigated by the achromatic amplification of the sample by foreground clusters).  Some of us have recently completed a large unbiased extragalactic submm survey (Scott et al.\\ 2002; Fox et al.\\ 2002; hereafter S02, F02) covering 260\\,arcmin$^2$ at 450 and 850\\,$\\mu$m. S02 detected 38 850-$\\mu$m sources at the $\\ge$\\,3.5\\,$\\sigma$ level ($N_{\\rm \\ge 8\\,mJy} = 320^{+80}_{-100}$\\,deg$^{-2}$) in the ELAIS N2 and Lockman Hole East regions.  This survey is very well-suited for determining the radio/submm spectral indices of SCUBA sources, and hence estimating the redshift distribution of the bright submm population. While the redshifts of individual sources are unlikely to be strongly constrained, the $N(z)$ can be determined statistically for a sufficiently large sample. At an 850-$\\mu$m detection threshold of $\\sim$\\,8\\,mJy, many sources will be detected by deep 1.4-GHz imaging. Moreover, any $\\sim$\\,8-mJy submm source {\\it not} detected at radio wavelengths can be ascribed a relatively robust and potentially exciting redshift constraint of $z\\ge 3$.  The redshifts of the more distant fraction can be constrained further using flux ratios that are more effective at $z\\rm\\gs 3$, e.g.\\ $S_{\\rm 850\\mu m}/S_{\\rm 1.25mm}$ (Eales et al.\\ 2002; see also Hughes et al.\\ 2002).  F02 presented shallow, $\\sim$\\,12$''$-resolution radio data from Ciliegi et al.\\ (1999) and de Ruiter et al.\\ (1997) for the 8-mJy survey regions.  With noise levels of $\\sim$\\,30\\,$\\mu$Jy\\,beam$^{-1}$, limits of $z\\gs 1$ could be set for most of the bright submm galaxy population. In the next section, we describe deep, high-resolution imaging ($\\sigma$ = 5--9\\,$\\mu$Jy\\,beam$^{-1}$, 1.4$''$ {\\sc fwhm}) of the 8-mJy survey regions. In \\S3 we use these maps to successfully identify robust radio counterparts for 60 per cent of the submm sources, and to refine the original submm sample via the excision of six sources which (in line with statistical expectation) appear to be the result of confusion. Next, in \\S4, we exploit the improved positional information provided by the 1.4-GHz maps to identify optical and/or near-IR host galaxies in new images, and to exclude possible counterparts where the radio data indicate blank fields ($V, R, I \\gs\\rm 26$, $K\\gs\\rm 21$). We go on to determine the redshift-sensitive submm-to-radio spectral indices for an unbiased sample of 30 sources from the $\\ge$\\,3.5\\,$\\sigma$ 8-mJy sample. Finally, in \\S5 we discuss the implications of the results of this multi-frequency follow-up study for the nature and redshift distribution of the luminous submm galaxy population.  Throughout we adopt a flat cosmology, with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$\\,km\\,${\\rm s^{-1}}$\\,Mpc$^{-1}$. ",
        "conclusions": "\\begin{enumerate} \\item We describe deep 1.4-GHz imaging of the 8-mJy survey regions in ELAIS N2 and Lockman East. These detect 60 per cent of the 30 submm-selected galaxies in our sample, enabling us to constrain the positions of these sources to better than 1$''$ and thereby identify host galaxies in other wavebands. \\item We present new optical and IR imaging and, based on the new positional information from the radio map, we find robust counterparts to 90 per cent of the radio-detected galaxies. Identifications based on colour are made for several more. \\item At least 60 per cent of the radio-detected optical/IR host galaxies display  highly-structured or distorted morphologies, suggestive of merging or interacting systems. \\item Almost one half of the optical/IR host galaxies are found to contain very or extremely red components. In addition, as many as ten of the optical/IR counterparts are composite systems comprising blue and red components separated by a few arcsec (tens of kpc at the relevant redshifts).  The strong internal colour gradients within these systems may be indicative of patchy dust obscuration. \\item Contrary to popular belief, virtually all of the host galaxies to the radio-detected population are sufficiently bright to justify spectroscopic observations with 8-m telescopes. We caution that redshifts require confirmation via CO detections before optical/IR host galaxies can be considered robust associations. \\item {\\it XMM-Newton} X-ray data for Lockman are presented, as well as {\\it Chandra} data for ELAIS N2. We detect four submm-selected galaxies, only one of which would have been identified as an AGN via its radio characteristics. \\item The diversity of the submm galaxy population is highlighted. We identify a beguiling mixture of sources, including eight EROs (one associated with the lobe of a radio galaxy) and two sources with flat-spectrum radio emission. \\item We find that less than a quarter of the sample would have been recovered by targeting optically faint radio sources, underlining the selective nature of such surveys. \\item We exploit the radio/far-IR correlation using our well-matched radio and submm data, finding a {\\it conservative} lower limit of $<\\!z\\!> \\geq$\\,2.0 for the median redshift of bright submm-selected galaxies, or $<\\!z\\!> \\geq$\\,2.4 using spectral templates more representative of known submm galaxies. \\item We find tentative evidence for luminosity evolution, with the brightest sources ($\\ge$\\,8\\,mJy) tending to be the most distant. \\item Employing our estimated redshift distribution, we find that submm galaxies with $S_{\\rm 850\\mu m}\\sim\\rm 8$\\,mJy play an important role in cosmic star-formation history. They are responsible for a higher star-formation-rate density at $z$\\,$\\sim$\\,1--4 than the entire galactic zoo manages at $z$\\,$\\sim$\\,0, and for a similar density as the $z$\\,$\\sim$\\,3--4 LBG population when extrapolated to $S_{\\rm 850\\mu m}>\\rm 1$\\,mJy. \\end{enumerate}"
    },
    "0206/astro-ph0206118_arXiv.txt": {
        "abstract": "We present results on the spectroscopic study of the ionized gas in the high redshift radio galaxy USS0828+193 at $z=$2.57. Thanks to the high S/N of the emission lines in the Keck spectrum, we have been able to perform a detailed kinematic study by means of the spectral decomposition of the emission line profiles. This study reveals the existence of two types of material in this object: a) a low surface brightness component with   apparent quiescent kinematics consistent with gravitational motions and b) a  perturbed component with   rather extreme kinematics. The quiescent halo extends across the entire object for $\\sim$80 kpc. It is enriched with heavy elements and apparently ionized by the continuum from the active nucleus.  The properties of the quiescent halo and its origin are discussed in this paper. We propose that it could be part of a structure that surrounds the entire object, although its nature is not clear (a rotating disc? low surface brightness satellites? a cooling flow nebula? material ejected in galactic winds? other?). ",
        "introduction": "Extended Ly$\\alpha$ regions are a common feature of high redshift radio galaxies ($z>$2, HzRG) and quasars (Heckman et al. 1991; see also narrow band Ly$\\alpha$ images of HzRG in, e.g, Kurk et al. 2001; Chambers, Miley \\& van Breugel  1990; McCarthy et al. 1990b). Most morphological and kinematic studies are based on the high surface brightness regions. These regions are clumpy, irregular and often aligned with the radio axis. They are characterized by extreme kinema\\-tics, with measured FWHM and velocity shifts  $\\geq$1000 km s$^{-1}$ (Baum \\& McCarthy 2000,  Villar-Mart\\'\\i n, Binette \\& Fosbury 1999, McCarthy, Baum \\& Spinrad 1996), compared to values of $\\sim$few hundreds in low redshift radio galaxies (Baum, Heckman \\& van Breugel 1990, Tadhunter, Fosbury \\& Quinn 1989). Although gravitational motions cannot be rejected, it is likely that a perturbing mechanism is responsible for the extreme kinematics.  The apparent connection between the radio and the kinematic properties  suggests that such a mechanism could be shocks generated by the interac\\-tion between the radio outflow and the gas in situ (van Ojik et al. 1997). Galactic winds (Heckman, Armus \\& Miley 1990) are an alternative possibility.   In addition to these regions, low surface  brightness Ly$\\alpha$  halos  extending beyond the radio structures have been detected in some HzRG. The kinematic properties of such halos have been studied in detail only in one case, the radio galaxy MRC1243+036 (van Ojik et al. 1996). The halo shows quies\\-cent kinematics [FWHM(Ly$\\alpha$)$\\sim$ 250 km s$^{-1}$] compared to the  regions inside the radio structures [FWHM(Ly$\\alpha$)$\\sim$1200 km s$^{-1}$]. The authors propose that the halo is a large rota\\-ting gaseous disc originating from the accretion associated with the formation of the galaxy.   Such quiescent low surface brightness halos (LSBHs, hereafter) are important since they show the gas properties unaffected by  kinematic perturbations. We are undertaking a research program using high S/N Keck spectroscopy of a sample of high redshift radio galaxies ($z\\geq$2.5) whose goal is to search for kinematically unperturbed Ly$\\alpha$ halos in HzRG. We will study the kinematic and  ionization properties, as well as observed properties such as surface brightness, size and luminosity. Constraints  will also be set on the halo mass  and its possible origin.   The results on  the radio galaxy USS0828+193 are presented in this paper. In a forthcoming paper we will discuss the results on the rest of the sample. ",
        "conclusions": "By means of the spectral decomposition of the emission line profiles in USS0828+193, we have isolated the emission from a giant ($\\sim$80 kpc) reservoir of apparently kinematically unperturbed gas, whose kinematics is consistent with being gravitational in nature. Emission lines other than Ly$\\alpha$ (CIV, HeII) have been detected for the first time in such a halo. This implies that the gas is ionized (probably by the continuum from the AGN) and enriched with heavy elements. We find marginal evidence for continuum detection from the halo.  We propose that the low surface brightness halo  in USS0828+193 (and other HzRG) could be part of a gaseous reservoir that surrounds the entire object. We have consi\\-dered several possible scenarios to explain the origin of the halo. The gas could have settled in a rotating disc. This is suggested by the velocity pattern of the  halo emission, although 3-dimensional spectroscopy would be necessary to set tighter constrains on the kinematic patterns.  Evidence for giant gaseous discs has been found by other authors in low redshift radio galaxies and elliptical galaxies.  The halo could also be a group of individual clouds (galactic satellites?) in a virialized system. In this case, the expected mass is $\\sim$8$\\times$10$^{11}$ M$_{\\odot}$. A similar value is estimated if the halo is infalling towards the center of the potential well. Such values are similar to estimates of masses of low redshift radio galaxies, and smaller than typical masses of  cD elliptical galaxies at low redshift. If this result is confirmed, we conclude that USS0828+193 will not become like the cD galaxies we see at the present epoch, unless mergers are involved.  Alternative scenarios that we have discussed (also consistent with the observations) are a cooling flow from a hot phase and outflows (galactic winds). In this last case, star forming rates $\\sim$500 M$_{\\odot}$ yr$^{-1}$ are required, consistent with  estimates for Lyman break galaxies.   The properties of the low surface brightness Ly$\\alpha$ halo detected in the radio galaxy USS0828+193  are consistent  with the expectations for the original gaseous reservoir  from which the galaxy started to form. This does not necessarily imply that the galaxy is in the process of formation.  We have also discussed the possible link of the low surface brightness halos with a) the very extended envelopes found around galaxies of  different  morphological types and luminosities based on studies of absorption line systems in the spectra of background quasars b)  the giant Ly$\\alpha$ halos associated with some Ly break galaxies c) the very extended reservoirs (claimed to be discs) of neutral gas associated with some nearby radio galaxies and elliptical galaxies."
    },
    "0206/astro-ph0206097_arXiv.txt": {
        "abstract": " ",
        "introduction": "It was proposed theoretically a long time ago that a merger of a pair of equal-mass spirals gives rise to an elliptical galaxy (Toomre 1977). This has been confirmed observationally by near-IR studies of a few mergers which indeed show an $r^{1/4}$ de Vaucouleur's radial profile typical of ellipticals (Schweizer 1982, Joseph \\& Wright 1985, Wright et al. 1990, Stanford \\& Bushouse 1991 (SB91 hereafter)). Theoretical work, involving mergers of galaxies (simulated using stars alone), has shown that under some conditions mergers indeed give rise to ellipticals with the above radial profile (Barnes \\& Hernquist 1991, Hernquist 1992). The mergers simulated with star formation in gas-rich spirals (Bekki 1998a) also give a similar profile.  Despite this overall agreement, there are many issues which need to be answered. First, the origin of the $r^{1/4}$ profile is not physically well understood, although violent relaxation is believed to be responsible for the elliptical light distribution (Lynden-Bell 1967). Both collisionless relaxation in an isolated galaxy (van Albada 1982) as well as in a merger (e.g. Barnes 1988) are known to result in an $r^{1/4}$ profile. Second, it is not clear if an $r^{1/4}$ profile over the entire galaxy is the only outcome that is possible for a galaxy merger, because the parameter space for galaxy interaction has not yet been studied exhaustively by the N-body theoretical work in the literature. The real galaxy mergers will presumably tell us about the parameter space that needs to be explored in future theoretical work. Third, other properties such as asymmetry in morphology, twisted isophotes, boxy/disky isophotes, shells, etc. have been shown to be good indicators that an elliptical galaxy was formed by a merger event (Nieto \\& Bender 1989; Hernquist \\& Spergel 1992).  In order to shed light on some of these issues, in this paper we have analyzed the near infrared K$_s$ band images from the 2MASS database for a sample of galaxies that show evidence of an advanced merger and we have also deduced other structural parameters. The large size of the database allows us to construct a big sample. The K$_s$ band best traces the underlying mass distribution in a galaxy. Also, the extinction in the K$_s$ band is very low, which enables us to study the underlying stellar distribution free of the effects of dust obscuration and young star-forming regions. We have chosen a sample of colliding galaxies from the Arp (1966) and Arp-Madore (1987) catalogs based purely on their optical appearance e.g. galaxies that show strong signs of interaction such as tails and a distorted, puffed-up main body but have a single merged nucleus. A similar study was first done by SB91 for ten galaxies, except that they considered the profiles only along the major axis and the minor axis. The subsequent papers to look at this problem have chosen their samples from IRAS (James et al. 1999, Scoville et al. 2000, Zheng et al. 1999) and are therefore biased towards high-luminosity, distant galaxies. Our sample, on the other hand, is unbiased with respect to the IR luminosity of a galaxy.  In this paper, we extend the  study of SB91 for a much larger sample of interacting galaxies, and study their radial profiles as well as other structural parameters. We give new detections of the $r^{1/4}$ profile in four galaxies from the Arp catalogue and thus our work has extended the known number of mergers showing such profiles. We have looked at the full azimuthal data and thus the luminosity of the galaxies is sampled in a better way. We check what fraction of our sample shows the $r^{1/4}$ profile expected for an elliptical galaxy and compare our results with those of the SB91, and other previous work. We also study other structural parameters such as boxiness/diskiness of isophotes, wandering centres etc. that provide important additional information on mergers. In Section 2, we describe the criteria for selection for our sample and also the data acquired from the public sites. In Section 3, we describe the data analysis. In Section 4, we discuss the results and our conclusions are summarized in Section 5.  \\begin{table*} \\caption[]{The sample} \\label{sample} \\begin{tabular}{clllcc} \\hline Galaxy&Alt.name&RA(2000)&Dec(2000)&Type(RC3)&Type(SIMBAD)\\\\ \\hline &&&&&\\\\ Arp156&UGC5814&10:42:38.0&+77:29:41&-&-\\\\ Arp160&NGC4194&12:14:09.8&+54:31:39&IBm pec&I\\\\ Arp162&NGC3414 &0:51:16.3&+27:58:33&S0 pec&SB0\\\\ Arp163&NGC4670&12:45:17.0&+27:07:34&SB(s)0/a pec&S0\\\\ Arp165&NGC2418&07:36:37.9&+17:53:06&E&E\\\\ Arp186&NGC1614&04:33:59.9&-08:34:30&SB(s)c pec&I\\\\ Arp187&-&05:04:52.99&-10:14:51&-&-\\\\ Arp193&IC883 &13:20:35.3&+34:08:22&Im: pec&-\\\\ Arp212&NGC7625&23:20:30.8&+17:13:41&SA(rs)a pec&Sp\\\\ Arp214&NGC3718&11:32:35.7&+53:03:59&SB(s)a pec&SB0/Sa\\\\ Arp219&-&03:39:53.26&-02:06:47&-&SB+cG\\\\ Arp221&-&09:36:28.03&-11:19:49&-&-\\\\ Arp222&NGC7727&23:39:53.8&-12:17:36&SAB(s)a pec&Sa\\\\ Arp224&NGC3921&11:51:06.1&+55:04:39&(R$^\\prime$)SA(s)0/a pec&S0\\\\ Arp225&NGC2655&08:55:38.7&+78:13:28&SAB(s)0/a&S0/Sa\\\\ Arp226&NGC7252 &22:20:44.9&-24:40:41&(R)SA(r)&-\\\\ Arp230&IC51&00:46:24.3&-13:26:34&S0 pec?&-\\\\ Arp231&PGC2616&00:43:54.3&-04:14:43&-&S0+\\\\ Arp243&NGC2623&08:38:24.2&+25:45:01&-&-\\\\ AM0318-230&PGC12526&03:20:40.1&-22:55:50&S?&Sa\\\\ AM0337-312&NGC1406&03:39:22.6&-31:19:19&SB(s)bc: sp&Sc\\\\ AM1315-263&NGC5061 &13:18:04.9&-26:50:10&E0&E\\\\ AM2146-350&NGC7135&21:49:43.7&-34:52:42&SA0- pec&S\\\\ AM0501-632&PGC16567&05:01:30.1&-63:17:34&S0+? pec&S\\\\ AM1300-233&MGC-4-31-23&13:02:52.07 &-23:55:19&-&Ir\\\\ AM1306-282&PGC45596&13:09:10.4&-28:38:23&SAB(rs)s pec?&Sa\\\\ AM1324-431&PGC47188 &13:27:51.0&-43:25:50&S?&Sc\\\\ \\end{tabular} \\end{table*}  \\begin{table*} \\caption[]{Basic data} \\label{bdata} \\begin{tabular}{clllll} \\hline Galaxy&f$_{100}/$f$_{60}$&L$_{IR}$&M$_B$&radius\\arcsec\\ &K$_s$\\\\ & &($\\times$ $10^{10} L_{\\odot}$)&&2MASS&2MASS\\\\ \\hline &&&&&\\\\ Arp 156&2.61&$<$0.21&-17.142&26.8&10.534\\\\ Arp 160&1.21&7.62&-20.322&25.6&9.680\\\\ Arp 162&$<$4.8&$<$0.07&-19.606&64.4&8.059\\\\ Arp 163&1.69&$<$0.21&-17.9&30&10.612\\\\ Arp 165&-&$<$0.27&-20.76&40.6&9.037\\\\ Arp 186&1.01&39.4&-21.193&24.2&9.565\\\\ Arp 187&-&-&-&16.8&11.314\\\\ Arp 193&1.63&36.9&-20.38&19.6&10.895\\\\ Arp 212&1.90&1.45&-19.081&42.8&8.929\\\\ Arp 214&3.28&0.07&-19.818&80&7.852\\\\ Arp 219&2.02&$<$5.69&-21.411&19.2&11.462\\\\ Arp 221&1.63&-&-20.376&33.8&9.491\\\\ Arp 222&-&-&-20.546&68.6&7.797\\\\ Arp 224&$<$2.45&$<$2.42&-21.58&27.8&9.939\\\\ Arp 225&2.98&0.25&-20.803&80&7.223\\\\ Arp 226&1.76&5.31&-20.926&33.2&9.399\\\\ Arp 230&2.12&0.41&-18.089&25.6&10.364\\\\ Arp 231&-&-&-19.713&26.4&9.887\\\\ Arp 243&1.20&29.33&-21.205&19.8&10.505\\\\ AM0318-230&-&-&-20.568&17.2&11.166\\\\ AM0337-312&2.26&0.82&-18.506&80&8.823\\\\ AM1315-263&-&$<$0.07&-21.118&71.8&7.372\\\\ AM2146-350&3.5&$<$0.20&-19.579&51&8.998\\\\ AM0501-632&2.60&0.06&-17.111&24.2&10.851\\\\ AM1300-233&1.11&23.61&-21.282&37.4&10.281\\\\ AM1306-282&2.05&$<$0.12&-18.372&23.4&11.291\\\\ AM1324-431&1.31&$<$10.74&-21.293&18&12.239\\\\ \\end{tabular} \\end{table*}   \\begin{table*} \\caption{Properties of the sample galaxies} \\label{props} \\begin{tabular}{cccccccl} \\hline Galaxy&D&1\\arcsec & R$_e$$,$R$_d$&$\\mu$&C$_{31}$&$\\sigma$&Spectral type\\\\ &(Mpc)&(pc)&(kpc)&(mag arcsec$^{-2}$)&&(km sec$^{-1}$)&\\\\ \\hline &&&&&&&\\\\ Class I&&&R$_e$&$\\mu_e$&&&\\\\ &&&&&&&\\\\ Arp 156&25&121&1.03$\\pm$ 0.14&18.08&4.35&-&-\\\\ Arp 165&67&325&5.92$\\pm$ 0.76&18.37&4.98&268&-early type \\\\ Arp 193&93&451&3.89$\\pm$ 0.38&17.89&3.68&206&LINER, starburst\\\\ Arp 221&-&-&-&18.32&4.18&-&-\\\\ Arp 222&25&125&2.83$\\pm$ 0.46&17.43&5.23&200&-\\\\ Arp 225&18&91&3.26$\\pm$ 0.38&17.70&4.86&164&LINER,early type\\\\ Arp 226&64&312&3.28$\\pm$ 0.77&17.43&4.74&177&-\\\\ Arp 231&76&368&3.80$\\pm$ 0.45&17.77&4.13&-&-\\\\ Arp 243&72&351&1.58$\\pm$ 0.36&16.68&4.07&95&LINER, starburst\\\\ &&&&&&&\\\\ Class II&&&R$_d$&$\\mu_0$&&&\\\\ &&&&&&&\\\\ Arp 162&17&87&1.56$\\pm$ 0.04&16.37&5.50&250&LINER\\\\ Arp 186&63&308&1.69$\\pm$ 0.09&15.33&4.62&150?&AGN\\\\ Arp 187&164&795&4.05$\\pm$0.19&16.44&3.24&-&-\\\\ Arp 212&21&104&0.85$\\pm$ 0.03&15.26&3.33&-&starburst\\\\ Arp 214&13&65&1.72$\\pm$0.11&16.96&5.41&178&LINER\\\\ Arp 219&139&673&4.86$\\pm$ 0.22&17.01&3.47&-&-\\\\ Arp 224&79&383&2.91$\\pm$ 0.21&16.57&5.14&195&post starburst\\\\ Arp 230&23&113&0.75$\\pm$ 0.02&16.21&3.17&-&-\\\\ AM0318-230&142&692&4.76$\\pm$ 0.51 &17.34&3.88&-&-\\\\ AM0337-312&14&69&2.28$\\pm$ 0.07&16.21&3.33&-&-\\\\ AM1324-431&139&675&4.39$\\pm$0.34&17.60&3.83&-&-\\\\ AM1315-263&24&116&1.90$\\pm$ 0.04&15.65&5.66&194&-\\\\ AM2146-350&36&174&3.22$\\pm$ 0.23&17.17&5.22&-&LINER\\\\ &&&&&&&\\\\ No fit&&&&&&&\\\\ &&&&&&&\\\\ AM0501-632&14&69&0.41$\\pm$0.003&16.19&2.69&-&-\\\\ Arp 160&34&168&-&-&3.21&104&starburst,AGN\\\\ Arp 163&16&78&-&-&3.52&-&Wolf Rayet, HII\\\\ AM1300-233&86&417&-&-&4.47&-&Seyfert\\\\ AM1306-282&19&92&-&-&2.63&-&-\\\\ &&&&&&&\\\\ \\end{tabular} \\end{table*} ",
        "conclusions": "A study of a sample of morphologically selected disturbed galaxies with indications of mergers shows that in spite of their similar optical appearances, the luminosity and hence mass profiles of these galaxies fall into two distinct classes. Thus, the dynamical evolution is distinctly different in the two classes. Class I is made up of galaxies that are fit by an $r^{1/4}$ law and hence are completely relaxed systems. Class II objects show the presence of an exponential disk in their outer regions. A few objects cannot be fit by either law. Thus, the dynamical properties of these classes are distinct. These classes do not show any differences in various properties, such as IR luminosities, molecular gas content or the IR colours. We find that our sample of morphologically disturbed galaxies may be merger products in spite of not being bright in the IR. Class I objects are likely to have formed as a result of merging of equal mass galaxies and have relaxed completely. Class II objects are likely to be the result of unequal mass mergers in which the exponential disk of the primary progenitor has survived. The unrelaxed objects are likely to be recent mergers.  We have also studied the structural parameters such as B4 and the location of the centre of mass, that provides important information on mergers, in addition to the mass distribution as indicated by the luminosity profiles. It is seen that the concentration index C$_{31}$ for Class I objects shows a maximum between 4 and 5. For Class II objects, there is a clear demarcation between the active and the non-active galaxies. The active galaxies show a peak between 5 and 5.5 while the non-active galaxies show a peak between 3 and 3.5. Thus, it is seen that active galaxies with an exponential disk are more centrally concentrated than completely relaxed systems which in turn are more centrally concentrated than non-active galaxies with an outer exponential disk.  In summary, 13 out of 27 disturbed galaxies with indications of mergers are found to exhibit exponential disk behaviour. Hence, the existence of this class is a robust result and this class of galaxies requires further dynamical study."
    },
    "0206/astro-ph0206005_arXiv.txt": {
        "abstract": "We study the evolution of stars that may be the progenitors of common (long-soft) GRBs. Bare rotating helium stars, presumed to have lost their envelopes due to winds or companions, are followed from central helium ignition to iron core collapse. Including realistic estimates of angular momentum transport \\citep{HLW00} by non-magnetic processes and mass loss, one is still able to create a collapsed object at the end with sufficient angular momentum to form a centrifugally supported disk, i.e., to drive a collapsar engine.  However, inclusion of current estimates of magnetic torques \\citep{Spr01} results in too little angular momentum for collapsars. ",
        "introduction": "One of the most promising models for the ``long variety'' of gamma-ray bursts (GRBs) is the so-called \\emph{collapsar} model \\citep{Woo93}. It assumes that a sufficiently massive stellar core collapses into a black hole and the infalling outer layers form a disk around it. Energy dissipated in the disk or the rotation of the black hole itself is assumed to power a jet of high Lorentz factor ($\\Gamma\\sim200$) that escapes from the engine along the polar axis to large distance ($\\sim\\Ep{15}\\,\\cm$) and powers a GRB by interaction with the circumstellar medium or by internal shocks.  The traversal time for the relativistic jet through the hydrogen envelopes of typical massive stars is hundreds to thousands of seconds.  Thus, at the time of the GRB, bare helium stars, which have radii of only a few light seconds (about a solar radius), are required if the lifetime of the engine and the GRB are not to be short compared to the time it takes the jet to drill through the star.  Two essential ingredients for the collapsar model are a sufficiently massive core to form a black hole and a sufficient rotation rate at the time of collapse to allow the formation of a disk.  The question we address here is: \\emph{What can be expected for the rotation rates of massive stellar cores when they collapse?} ",
        "conclusions": "A bare helium star of low metallicity can retain enough angular momentum to form a centrifugally supported disk around a central black hole of $\\sim3\\,\\Msun$, as required by the collapsar model for GRBs. Without magnetic fields, the angular momentum is sufficient to form a Kerr black hole and support most or all of the star in an accretion disk.  However, if we include an approximate treatment of angular momentum transport by magnetic fields, the resulting spin rates become too low to form centrifugally supported disks in the inner part of the core. Even a binary helium star merger at the end of central helium burning might not be able to avoid this fate.  Mass loss can lead to an additional significant spin-down of the core, especially if magnetic fields couple it effectively to the envelope.  Even in case of Keplerian surface rotation, the core rotation needs to decouple before carbon ignition in order to make a Kerr black hole.  The dynamo process recently proposed by \\citet{Spr01} seems too efficient to form collapsar progenitors from single stars or helium star mergers.  This is even more so for the magnetic field modeling suggested by \\citet{SP98}.   \\begin{theacknowledgments} We thank Henk Spruit for a preview of his work and many helpful discussions.  This work has been supported by the NSF (AST-9731569), NASA (NAG5-8128), the DOE (B347885), and the AvH (FLF-1065004). \\end{theacknowledgments} \\vspace{-0.1in}"
    },
    "0206/astro-ph0206233_arXiv.txt": {
        "abstract": "{ We use ultraviolet space-based ({\\it FUSE, HST}) and optical/IR ground-based (2.3m MSSSO, NTT) spectroscopy to determine the physical parameters of six WC4-type Wolf-Rayet stars in the Large Magellanic Cloud. Stellar parameters are revised significantly relative to Gr\\\"{a}fener et al. (1998) based on improved observations and more sophisticated model atmosphere codes, which account for line blanketing and clumping. We find that stellar luminosities are revised upwards by up to 0.4 dex, with surface abundances spanning a lower range of 0.1$\\le$ C/He $\\le$ 0.35 (20--45\\% carbon by mass) and O/He$\\le$0.06 ($\\leq$10\\% oxygen by mass). \\\\ Relative to Galactic WC5--8 stars at known distance, and analysed in a similar manner, LMC WC4 stars possess systematically higher stellar luminosities, $\\sim$0.2\\,dex lower wind densities, yet a similar range of surface chemistries. We illustrate how the classification C\\,{\\sc iii} $\\lambda$5696 line is extremely sensitive to wind density, such that this is the principal difference between the subtype distribution of LMC and Galactic early-type WC stars.  Temperature differences do play a role, but carbon abundance does not affect WC spectral types. We illustrate the effect of varying temperature and mass-loss rate on the WC spectral type for HD\\,32257 (WC4, LMC) and HD\\,156385 (WC7, Galaxy) which possess similar abundances and luminosities.\\\\ Using the latest evolutionary models, pre-supernova stellar masses in the range 11--19 $M_{\\odot}$ are anticipated for LMC WC4 stars, with 7--14 $M_{\\odot}$ for Galactic WC stars with known distances. These values are consistent with pre-cursors of bright Type-Ic supernovae such as SN\\,1998bw (alias GRB 980425) for which a minimum total mass of C and O of 14$M_{\\odot}$ has been independently derived. ",
        "introduction": "The ultimate fate of the most massive stars is likely to be a Type Ib or Ic Supernova (SN) explosion. Type Ib's, conspicuous for the absence of hydrogen in their spectra, likely correspond to WN-type Wolf-Rayet (WR) stars, whilst WC or WO stars are thought to be responsible for Type Ic SN, since both hydrogen and helium are absent  from their (early) spectra. Are  the properties of WC stars immediately prior to their proposed SN explosion consistent with that of such SN? Models of WR stars and SN are now sufficiently advanced to facilitate such comparisons for the first time.   Over the past couple of decades, spectroscopic tools for the quantitative analysis of WR stars have advanced sufficiently (e.g. Hillier \\& Miller 1998) to permit the reliable determination of abundances (Herald et al. 2001), masses (De Marco et al. 2000) and ionizing fluxes (Crowther et al. 1999).  Armed with these tools, we are now in an  unprecedented position to investigate such stars in the Local Group and beyond (e.g. Smartt et al. 2001). Already, high quality observations spanning UV to IR wavelengths, are possible for individual stars in our Galaxy or the Magellanic Clouds with new instruments, such as the {\\it Far-Ultraviolet Spectroscopic Explorer} ({\\it FUSE}, Moos et al. 2000).  WR stars in the LMC have been the focus of several spectroscopic investigations. WN stars have been studied by Conti \\& Massey (1989), whilst Hamann \\& Koesterke (1998), and references therein, have determined their  quantitative properties, for which a negligible metallicity effect was remarked upon relative to Galactic counterparts (the heavy metal content of the LMC is $\\sim 0.4Z_{\\odot}$, Dufour 1984). Torres et al. (1986), Smith et al. (1990) and Barzakos et al. (2001) have compared LMC WC stars, with Gr\\\"{a}fener et al. (1998) presenting studies of six WC4 stars. In the latter work, WC4 stars were found to possess remarkably uniform (and high) carbon and oxygen abundances, plus a wide range of stellar luminosities. Results were inconclusive whether the mass-loss rates of LMC WC stars are similar to, or lower than, Galactic counterparts. We return to the study of these stars in this paper, since we possess improved and more extensive spectroscopy, including far-UV and near-IR datasets, plus better modelling tools (Hillier \\& Miller 1998).  \\begin{table} \\caption{Observing log for LMC WC4 stars, including narrow-band photometry from Torres-Dodgen \\& Massey (1988), and catalogue numbers from Breysacher (1981, Br) and Breysacher et al. (1999, BAT).} \\label{table1} \\begin{tabular}{l@{\\hspace{1mm}}r@{\\hspace{1mm}}r@{\\hspace{1mm}}c@{\\hspace{1mm}} r@{\\hspace{1mm}}r@{\\hspace{1mm}}r@{\\hspace{1mm}}r@{\\hspace{1mm}}r}\\\\ \\hline HD     &    Br  & BAT &  $v$   &  $b-v$ & {\\it FUSE} & {\\it HST} & MSSSO & NTT \\\\ \\hline 32125  &  7   &  9    & 15.02  & 0.10    & --  & Sep 94 &Dec 97 & -- \\\\ 32257  &  8   &  8    & 14.89  & 0.13    & --  & Nov 94 &Dec 97 & -- \\\\ 32402  & 10   &  11   & 13.89  & $-$0.06 &Nov 01  & Apr 95 &Dec 97 & Sep 99\\\\ 37026  & 43   &  52   & 14.04  & $-$0.03 &  Feb 00  & Nov 94 &Dec 97 & -- \\\\ 37680  & 50   &  61   & 14.03  & $-$0.01 &  Feb 00 & Nov 94 &Dec 97 & -- \\\\ 269888 & 74   & 90    & 15.41  & 0.19    & --  & Jun 95 &Dec 97 & -- \\\\ \\hline \\end{tabular} \\end{table}   We present our new observations in Sect.~2, and discuss the present set of model calculations in Sect.~3. Individual results are presented in Sect.~4 and compared with the previous study of Gr\\\"{a}fener et al. Quantitative comparisons are made with Galactic counterparts that have been analysed in a similar manner (Dessart et al. 2000; Hillier \\& Miller 1999) in Sect.~5, together with evolutionary expectations. Finally, the possibility that WC stars provide the precursors to Type Ic SN are discussed in Sect.~6, via the comparison of accurate WC masses with those of the CO-cores determined for recent SN Ic explosions (e.g. Iwamoto et al. 1998, 2000).   \\begin{figure*}[ht!] \\vspace{10cm} \\special{psfile=crowther1.eps hoffset=+20 voffset=-225 hscale=70 vscale=70} \\caption{{\\it Panels (a-c)}: Synthetic spectra for selected C\\,{\\sc iv} line profiles with parameters $T_{\\star}=90$kK, $\\log L/L_{\\odot}$=5.70, $v_{\\infty}$=3000 km\\,s$^{-1}$, C/He=0.14, O/He=0.02 by number. An unclumped model is shown in solid ($f$=1, with $\\dot{M} = 1.1 \\times 10^{-4} M_{\\odot}$ yr$^{-1})$, moderately clumped as a dotted line ($f$=0.1, with $\\dot{M} = 3.5 \\times 10^{-5} M_{\\odot}$ yr$^{-1})$), and highly clumped as a dashed-line ($f=0.01$, with $\\dot{M} = 1.1 \\times 10^{-5} M_{\\odot}$ yr$^{-1}$. Electron scattering wings, and line intensities for the case of $\\lambda\\lambda$5801--12, are very sensitive to clumping, as is the shape and location of $\\lambda\\lambda$1548--51. {\\it Panels (d-f):} Comparison of moderately clumped ($f$=0.1) models with observations of HD\\,32402, showing generally good agreement. Ordinate units are erg\\,cm$^{-2}$\\,s$^{-1}$\\,\\AA$^{-1}$.} \\label{clump} \\end{figure*} ",
        "conclusions": "We have used far-UV, UV, optical and near-IR spectroscopy together with line blanketed model atmospheres to re-evaluate the stellar parameters of six LMC WC4 stars. We derive stellar luminosities which are a factor of $\\sim$2 higher than previously established for these stars (Gr\\\"{a}fener et al. 1998), and systematically higher than Galactic counterparts at known distance. Mass-loss rates are lower than previously determined due to the inclusion of clumping in our models. Several lines of evidence, in both WN and WC stars, strongly suggest that the winds of WR are strongly clumped.  We find that the LMC WC4 stars possess very similar surface abundances to Galactic WC5--8 stars, although their wind densities are systematically lower, by $\\sim$0.2 dex. We propose that the lower heavy metal content of the LMC is responsible for their lower mass-loss rates via their (still unproven) radiatively driven winds. A metallicity dependence of $\\sim Z^{0.5}$ would produce this weak effect, comparable to that predicted for O stars as a function of metallicity ($\\dot{M} \\propto Z^{0.5-0.7}$: Kudritzki \\& Puls 2000; Vink et al. 2001). This relatively minor difference quantitatively explains their different subtype distribution (recall Fig.\\ref{wc4_wc7}). Temperature appears to be a secondary criteria in distinguishing WC4--WC7 stars. Late-type WC stars, especially WC9, do appear to be rather different, with systematically lower stellar temperatures. Looking at the broader picture of the subtype  distribution in other Local Group galaxies, other factors clearly come into play.  If metallicity were the sole factor which determined the wind strengths of WC stars, those galaxies with metallicities close to Galactic (e.g. M33 and M31) would have a large, late-type WC population, which is apparently not the case (Massey \\& Johnson 1998), unless these are all obscured in the visible by circumstellar dust. Recent evidence suggests that most (or all) WC9 stars are close binaries with time-dependent dust production from interactions between the winds of the WC and OB companion (e.g. Tuthill et al. 1999). Future quantitative studies of the properties of these and more distant WR stars will hopefully help address these issues. Regardless of these details, the identification of a heavy metallicity dependence of WR winds has a major impact on their relative contribution to the hard ionizing photons in young starburst galaxies (Smith et al. 2002).  The association between spectral type and wind density amongst LMC and Galactic WCE stars also has relevance to the long standing debate amongst [WC]-type central stars of Planetary Nebulae (CSPN). It has long been recognised that [WC] stars positively avoid WC5-7 spectral types, and either cluster around WC4/WO or WC8--10 (e.g. Fig.~2 in Acker et al. 1996; Crowther et al. 1998). A wind density origin for the differences between WC4--7 stars (rather than temperature) can therefore explain the observed clustering as due to reduced wind densities amongst [WC]-type CSPN relative to massive Galactic WC stars, without the need to infer rapid evolution, for instance. Clearly, one still would need to understand why [WC]-type CSPN possess weaker winds than their massive cousins, but the differences would not need to be great -- a factor of $\\sim$2 would suffice.  Finally, we determine current masses of LMC and Galactic WC stars including the (minor) effect of wind darkening (Heger \\& Langer 1996). Based on remaining lifetimes from evolutionary predictions we estimate final pre-SN masses of 11--19$M_{\\odot}$ for LMC WC stars, and 7--14$M_{\\odot}$ for Galactic WC stars with known distances. These values are consistent with WC stars being the immediate precursors to luminous Type-Ic supernova explosions, including SN\\,1998bw and SN\\,1997ef."
    },
    "0206/astro-ph0206143_arXiv.txt": {
        "abstract": "{ We have performed abundance analysis of two slowly rotating, late A-type stars, HD~32115 (HR~1613) and HD~37594 (HR~1940), based on obtained echelle spectra covering the spectral range 4000-9850 \\AA. These spectra allowed us to identify an extensive line list for 31 chemical elements, the most complete to date for A-type stars. Two approaches to abundance analysis were used, namely a ``manual'' (interactive) and a semi-automatic procedure for comparison of synthetic and observed spectra and equivalent widths. For some elements non-LTE (NLTE) calculations were carried out and the corresponding corrections have been applied. The abundance pattern of HD~32115 was found to be very close to the solar abundance pattern, and thus may be used as an abundance standard for chemical composition studies in middle and late A stars.  Further, its H$\\alpha$ line profile shows no core-to-wing anomaly like that found for cool Ap stars and therefore also may be used as a standard in comparative studies of the atmospheric structures of cool, slowly rotating Ap stars.  HD~37594 shows a metal deficiency at the level of -0.3 dex for most elements and triangle-like cores of spectral lines. This star most probably belongs to the $\\delta$ Sct group. ",
        "introduction": "Solar photospheric abundances are usually used as a reference for any abundance study. However, it has been shown that the mean abundances for some elements obtained for large homogeneous groups of stars can differ significantly from solar abundances. For example, Luck \\& Lambert (1981, 1985) and later Takeda \\& Takada-Hidai (1998) obtained an oxygen deficiency of 0.3 dex relative to the Sun for a large group of A to K supergiants. There are some indications that the sun may be slightly iron-rich (+0.09 dex on average) as compared with solar-like stars of the same age (Gustafsson 1998).  Extensive abundance studies of the superficially normal sharp-lined B- and F-type dwarfs were performed by S.\\ Adelman and co-workers (see Adelman et al.\\ 2000 and references therein). They found that many of the investigated stars show Am-type phenomena, and only 5 stars with effective temperatures in the range 6700--9000 K have abundances close to the solar photospheric values. Their study was performed mainly in the blue spectral region (up to 5000 \\AA), and therefore abundances of C, N, O, Na, S, K were absent or rather questionable. Varenne \\& Monier (1999) published abundances for 48 A- and F-type dwarfs in the Hyades open cluster. They have only two stars hotter than 6650 K with \\vsini\\ $\\leq$ 25 \\kms. Because of the limited spectral regions that were observed abundances for only 11 elements have been derived. Half of these abundances were based on 1 or 2 spectral lines.  Most sharp-lined stars with temperatures in the range 6700--8000 K belong to different groups of peculiar stars: Ap, roAp, Am, or $\\lambda$ Boo groups. They are usually believed to have the same or nearly the same atmospheric structure as normal stars. However, Cowley et al.\\ (2001) found a pronounced core-to-wing anomaly in the Balmer lines of some Ap stars. They could not compare the H$\\alpha$ line profiles of these stars with normal solar abundance stars because of the lack of the reliable spectroscopic standards in this temperature region. We therefore decided to perform an accurate spectroscopic investigation of sharp-lined late A- to early F-type stars which are classified as normal MS stars, using observations of a wide spectral region 4000--9850 \\AA. From tabulated rotational velocities (Abt \\& Morrell 1995) we extracted 27 stars classified as normal A3--F5-type stars of luminosity classes III-V with \\vsini$\\leq$ 25 \\kms. Among them only 13 stars have metallicities in the range $|[M]|\\leq$0.15 on the basis of temperature-gravity-metallicity calibrations of the observed Str\\\"{o}mgren colours (a procedure which will be described below). We did not consider the remainder of the stars, the majority of which have low metallicities $[M]<$-0.15. In this paper we present the results of careful atmospheric parameter and abundance determinations for two stars: HD~32115 and HD~37594. These stars were classified as A8IV (HD~32115) and A8Vs (HD~37594) by Cowley et al.\\ (1969). ",
        "conclusions": "The abundances found for HD~32115 and HD~37594 relative to the sun are shown in Fig.3 by filled and open circles respectively.  Within the typical error limits of $\\pm$0.15 dex HD~32115 is a solar abundance star (a mean metallicity $[M] = -0.04\\pm0.11$). Therefore this star may be used as a chemical standard in studies of cool peculiar stars. We also get a standard for further investigation of the hydrogen wing-to-core anomaly found in cool Ap stars by Cowley et al.\\ (2001). Fig.4 shows a comparison between H$\\alpha$ line profiles in HD~32115 and in one of the pulsating Ap (roAp) stars, HD~24712, with the same effective temperature. It is evident from this comparison that the Ap-star anomaly occurs in that part of the hydrogen line core which is not reproduced by our current models.  \\begin{figure}[th] \\includegraphics[width=88mm]{elemental_diff_7250.ps} \\caption{The observed relative abundances in HD~32115 (filled circles) and in HD~37594 (open circles).} \\label{fig3} \\end{figure}  \\begin{figure}[th] \\includegraphics[width=88mm]{h_alpha_comp_ap.ps} \\caption{A comparison between H$\\alpha$ line profiles observed in HD~32115 (thick full line) and in roAp star HD24712 (dashed line), and calculated for the adopted 7250-ODF model (thin full line).} \\label{fig4} \\end{figure}   HD~37594 is slightly metal-deficient. Its mean metallicity $[M]=-0.26\\pm0.16$ is close to the value $[M]=-0.15\\pm0.04$ obtained from photometric calibrations.  This error in the metallicity obtained from photometric calibrations is based only on the errors in the photometric indices. No systematic errors of the calibration itself are included, which may be up to 0.1. Further support for the high accuracy of the metallicities derived from the Str\\\"{o}mgren photometry for normal stars is provided by the results of abundance analysis of 28 And = HD~2628 (Adelman et al.\\ 2000) and $\\sigma$ Boo~A=HD~128167 (Adelman et al.\\ 1997). [M]=-0.16$\\pm$0.29 and -0.29$\\pm$0.37 were derived for these stars from abundances versus -0.13 and -0.35 from photometric calibration. The abundance pattern in HD~37594 is similar to that in 28 And (\\teff=7250 K, \\lgg=4.2, \\vt=2.3 \\kms, \\vsini=9 \\kms -- see Adelman et al.\\ 2000) which belongs to the $\\delta$ Sct group. The smaller scatter derived in the present paper is explained by more accurate abundance determinations with NLTE and hfs effects taken into account as well as more careful choice of the unblended lines which is possible in the red spectral region.  Abundances obtained for HD~32115 and HD~37594 do not follow the predicted abundance pattern for $Z\\leq$28 from consistent stellar evolution models calculated with radiative forces, opacities and diffusion (Turcotte et al.\\ 1998). Models for 1.5 $M/M_{\\sun}$ predict that elements with 5~$<$~$Z$~$<$~20 will be underabundant by 0.5 dex relative to Fe, and that iron-peak elements will be generally slightly overabundant even at early evolutionary phases on the MS. This was not observed in either of the stars analysed here.  To confirm the derived abundances we have used the software \\vwa\\ (Bruntt et al.\\ 2002) which was developed to be able to make fast semi-automatic abundance analysis. We find the same abundances as the more careful classical approach. At present \\vwa\\ is being applied to the study of the primary target candidate stars for the asteroseismology missions COROT and  R\\o mer (Bruntt et al.\\ 2002). Several of these stars have moderate or high \\vsini\\, and the amount of work needed to carry out the analysis using \\vwa\\ is reduced substantially."
    },
    "0206/astro-ph0206469_arXiv.txt": {
        "abstract": "Recently, high-resolution {\\sl Chandra} observations revealed the existence of very sharp features in the X-ray surface brightness and temperature maps of several clusters \\citep{vikhlinin01}.  These features, called {\\em ``cold fronts''}, are characterized by an increase in surface brightness by a factor $\\gtrsim 2$ over 10-50~kpc, accompanied by a {\\it drop} in temperature of a similar magnitude. The existence of such sharp gradients can be used to put interesting constraints on the physics of the intracluster medium (ICM), if their mechanism and longevity are well understood. Here, we present results of a search for cold fronts in high-resolution simulations of galaxy clusters in cold dark matter (CDM) models. We show that sharp gradients with properties similar to those of observed cold fronts naturally arise in cluster mergers when the shocks heat gas surrounding the merging sub-cluster, while its dense core remains relatively cold.  The compression induced by supersonic motions and shock heating during the merger enhance the amplitude of gas density and temperature gradients across the front. Our results indicate that cold fronts are non-equilibrium transient phenomena and can be observed for a period of less than a billion years.  We show that the velocity and density fields of gas surrounding the cold front can be very irregular which would complicate analyses aiming to put constraints on the physical conditions of the intracluster medium in the vicinity of the front. ",
        "introduction": "\\label{sec:intro}  Recently, high-resolution {\\sl Chandra} observations revealed the existence of very sharp features in the X-ray surface brightness and temperature maps of several clusters, including A2142 \\citep{markevitch00} and A3667 \\citep{vikhlinin01}. The sharp gradients in the \\xray surface brightness were first found in the {\\sl ROSAT} images of A2142 and A3667 along with signatures of substructures (i.e., several small groups and merging components embedded in a larger cluster).  Due to lack of spatially resolved temperature maps, these features were initially interpreted as shock fronts arising during cluster mergers \\citep{markevitch99}.  This interpretation turned out to be incorrect when detailed \\xray surface brightness and temperature maps were obtained with the {\\sl Chandra} satellite.  The maps showed that the increase in gas density and \\xray surface brightness by a factor of $\\gtrsim 2$ across the sharp features is accompanied by a {\\em decrease} in temperature of gas of a similar magnitude, the behavior opposite to that expected across shock fronts.  These sharp features were therefore named {\\em ``cold fronts''}.  \\citet{markevitch00} and \\citet{vikhlinin01} interpreted these fronts as a boundary of ``the dense subcluster core that has survived a merger and ram pressure stripping by the surrounding shock-heated gas.''  The existence of the sharp temperature gradients can be used to put interesting constraints on the conditions in the intracluster medium (ICM) in the vicinity of the front. In particular, the width of the observed front ($\\lesssim 5-10$~kpc) is several times smaller than the Coulomb mean free path for electrons in the ICM.  This indicates that the thermal conduction must be suppressed, at least across the front. In the case of A2142, \\citet{ettori00} find that the thermal conductivity has to be reduced by several orders of magnitude from the classical Spitzer value near the front.  The observed extent, $\\sim 500$~kpc, of the sharp boundary may also indicate that the Kelvin-Helmholtz instabilities are partially suppressed, most likely by magnetic fields parallel to the boundary.  The instabilities, expected to arise on the boundary if the hot gas flows along the front, would disrupt and widen the boundary.  Using a simple dynamical model that approximates the gas cloud as a dense spherical cloud, \\citet{vikhlinin01} find that the magnetic field of 7-16$\\mu$G parallel to the front must be present at the boundary to maintain its stability.  The detailed studies of cold fronts can, thus, provide new detailed insights into the physics of the ICM, including the efficiency of energy transport and the magnetic field strength in the ICM. The reliability of such models, however, depends on understanding the dynamics of gas motions during cluster mergers and, especially, in the vicinity of the cold fronts. The constraints on the transport mechanisms require also that we understand how long the cold fronts can survive dynamically.  In this paper we present results of the search and analysis of features similar to the observed cold fronts in very high-resolution cosmological N-body$+$gasdynamics simulations of clusters forming in Cold Dark Matter (CDM) models.  The main goals of this paper are 1) to test the interpretation of the cold fronts as boundaries between hot shock-heated gas and dense cold gas of the merging sub-clump; 2) identify the situations and mechanisms that can produce cold fronts and 3) analyze the dynamics and density fields of gas and dark matter in the vicinity of the front.  Although numerous extensive theoretical analyses of cluster mergers have been done in the past, the typical spatial resolution in studies of cluster formation in the full cosmological context is too low to match the superb spatial resolution of the {\\sl Chandra} images. Eulerian gasdynamics codes based on the modern high-resolution shock capturing techniques are capable of resolving sharp density and temperature gradients of arbitrary amplitude within $1-2$ grid cells and are thus well suited for studies of shocks and cold fronts in cluster mergers. However, high-dynamic range required in self-consistent simulations of cluster formation limited use of the Eulerian codes to controlled merger experiments \\citep[e.g.,][]{roettiger97a,quilis98,ricker01}.  The results of these simulations were widely used for physical interpretation of the new high-resolution \\xray observations of merging clusters.  More recently, the advent of cosmological codes using the adaptive mesh refinement (AMR) \\citep[e.g.,][]{bryan97,kravtsov02,teyssier02} allowed to achieve the dynamic range of $\\gtrsim 10^4-10^5$ with Eulerian gasdynamics algorithms, thus making possible high-resolution self-consistent cluster simulations in a realistic cosmological setting.  As we will show below, these simulations can resolve sharp gradients in density and temperature fields of the ICM on scales of $\\sim 10$~kpc, approaching the typical resolution of {\\sl Chandra} images.  The paper is organized as follows. In $\\S$2, we discuss the cluster simulations used in our analysis.  In $\\S$3, we discuss and illustrate the detailed structural evolution of the cluster gas and surrounding filaments during cluster mergers and accretion of matter.  In \\S~4 we present cold fronts identified in numerical simulations and discuss their origin, structure, dynamical properties, and detectability. We discuss our results and summarize our conclusions in \\S~5.   \\begin{figure*}[t] \\centerline{ \\epsfysize=1.7truein  \\epsffile{Plot_LowerRes/fig1a.ps} \\epsfysize=1.7truein  \\epsffile{Plot_LowerRes/fig1b.ps} \\epsfysize=1.7truein  \\epsffile{Plot_LowerRes/fig1c.ps} \\epsfysize=1.7truein  \\epsffile{Plot_LowerRes/fig1d.ps} }  \\centerline{ \\epsfysize=1.7truein  \\epsffile{Plot_LowerRes/fig1e.ps} \\epsfysize=1.7truein  \\epsffile{Plot_LowerRes/fig1f.ps} \\epsfysize=1.7truein  \\epsffile{Plot_LowerRes/fig1g.ps} \\epsfysize=1.7truein  \\epsffile{Plot_LowerRes/fig1h.ps} } \\caption{ \\footnotesize The maps of projected density (top) and entropy (defined as $Tn^{-2/3}$, bottom) of a simulated $\\Lambda$CDM cluster at four different redshifts in a $60h^{-1}$~kpc slice centered on the central density peak. The maps are color-coded on a $\\log_{10}$ scale in units of cm$^2$ (density) and keV cm$^2$ (entropy). The size of the region shown is 8$h^{-1}$Mpc (10\\% of the entire simulation volume). The entropy maps reveal a very complex entropy distribution of the gas. Both the filaments and the forming cluster are surrounded by strong accretion shocks.  Note, however, that the accretion shock around the cluster is very aspherical and does not penetrate into the filament; relatively low-entropy gas accreting onto cluster along the filament does not pass through the strong virial shocks and can be traced all the way to the central $0.5\\hMpc$ of the cluster. } \\label{clevol} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  We used high-resolution $N$-body$+$gasdynamics cluster simulations in CDM models to search for counterparts of {\\em ``cold fronts''} within cluster ICM recently discovered in {\\sl Chandra} observations.  The observed cold fronts are very sharp features in the X-ray surface brightness and temperature maps found in several clusters.  These features are characterized by an increase in surface brightness by a factor $\\gtrsim 2$ accompanied by a {\\it drop} in temperature of a similar magnitude over the scale of 10-50~kpc.  The existence of such sharp features can put interesting constraints on the properties of the intracluster medium (ICM), including the efficiency of energy transport \\citep{ettori00} and the magnetic field strength in the ICM \\citep{vikhlinin01}.  In this paper, we investigate whether such features can be produced naturally in clusters forming in hierarchical models and study their origin, structural properties and detectability.  We find that cold fronts appear to be common and often arise in major and minor cluster mergers.  In the preceding sections we discussed two specific cases of cold fronts found in the simulation of the $\\Lambda$CDM cluster. In the most spectacular case, the cold front arises in a major merger when the merging subclump, moving slightly supersonically, is undergoing tidal disruption (see distribution of DM and velocity field in Fig.~\\ref{dm_temp}). The motion of the gas, which appears to be in the process of being ``sloshed out'' of the potential well of its DM halo, drives the shock and leads to a steepening of density and temperature gradients which produces the cold front.  The surface brightness and temperature profiles across the front show that a drop of temperature over the scale of $20-30\\hkpc$ is accompanied by a jump in surface brightness by a factor of $\\gtrsim 2$, behavior very similar to the profiles of observed cold fronts. In addition, we find that similarly to observations pressure changes smoothly accross cold fronts in simulations. The spatial extent of the cold front in simulation ($\\sim 0.5\\hMpc$) is comparable to the extent of two prominent observed cold fronts in clusters A2142 and A3667 \\citep{markevitch00,vikhlinin01}. Note also that the cold front is preceded by a bow shock ($\\sim 200$~kpc from the cold front). In this case the jump in surface brightness by a factor of $\\gtrsim 5$ is accompanied by an {\\it increase} in temperature of a similar magnitude over the scale of $\\sim 40$~kpc. The shock front has a much lower \\xray surface brightness than the gas around the cold front and therefore would be difficult to observe. Nevertheless, indications of a similar bow shock preceding the cold front were found in A3667.  In another discussed case, the cold front arises in a minor merger when a merging subclump reaches the inner regions of the larger cluster and gas in front of the subclump is significantly compressed. The compression sharpens and enhances the amplitude of gas density and temperature gradients across the front. Here the subcluster is on the first approach to the main cluster's core and is not yet significantly disturbed by tides. The relatively cooler gas of the merging sub-clump is trailed by low-entropy ($\\sim$1-2~keV) intergalactic gas in the direction of the subclump's motion. Such elongated regions of low-temperature gas trailing a subclump, if observed, can therefore be identified with regions where filamentary material is accreting onto cluster and could potentially be targets for studies of properties of intergalactic material. Unfortunately, the \\xray surface brightness of this gas is very low and detecting such regions will be a challenge. The best chance of observing such gas in candidate regions would be X-ray and UV metal absorption lines in spectra of bright background sources \\citep[e.g.,][]{hellsten98}.  Cold fronts in both observed and simulated clusters seem to come in all shapes and sizes and are rather frequent, as reports of many new instances of cold fronts are published \\citep[e.g.][]{sun02,forman02,markevitch02}.  In the two clusters that we analyzed, we identified many ``small-scale'' cold fronts many of which are located in the outskirts of the main cluster progenitor. Most of these cold fronts would probably not be detected in distant clusters with typical exposures and field sizes of Chandra observations, although they are similar to those observed in the Fornax cluster \\citep{forman02}.  The fact that large-scale ($\\sim 0.5$~Mpc) cold fronts arise in major mergers and survive for approximately dynamical time ($< 1$~Gyr) implies that their frequency will be similar to the frequency of major mergers. The latter is a very strong function of redshift ($\\propto (1+z)^{\\sim 3}$, \\citep[see, e.g.][]{gottlober01}) and may not be very high at $z=0$ (depending on cosmology and the power spectrum normalization), but be increasingly larger for higher redshift clusters.  Indeed, the large-scale cold front that we discussed occurs at $z\\approx 0.4$. Observationally, cluster ZW3146 at $z=0.29$ exhibits several cold front like features \\citep{forman02}. More recently, \\citet{markevitch02} reported results of a systematic search for cold fronts in archival {\\sl Chandra} observations of 30 apparently relaxed clusters. Surprisingly, they find cold-front like features in the large fraction of the clusters.  Larger samples of both observed and simulated clusters will be useful to determine the frequency of cold fronts.  Our results support the interpretation of cold fronts as boundaries between hot ICM of the main cluster and ``the dense subcluster core that has survived a merger and ram pressure stripping by the surrounding shock-heated gas'' \\citep{markevitch00}. Based on this interpretation, \\citet{vikhlinin01} proposed the detailed dynamical model of the cold front observed in cluster A3667 assuming a spherical dense gas subclump in a spherical dark matter potential moving through the uniform ICM. This model was used to place constraints on the magnetic field configuration and strength and plasma conductivity in the vicinity of cold front. The conductivity can be constrained simply by the sharpness of the observed front if a certain lifetime is assumed for the front. The constraint on the magnetic field configuration, on the other hand, is obtained from the fact that no indication of Kelvin-Helmholtz instabilities is present in observations.  Such instabilities are expected to develop quickly for a regular high-velocity flow of hot surrounding gas around the dense subclump along the cold front boundary. The power of these constraints relies on the validity of the underlying assumptions. Of these, the assumption of regular flow is particularly worrisome, given the complicated flow structure apparent from entropy and temperature maps shown in Figs.~\\ref{clevol} and~\\ref{xsbt_map}.  Indeed, Figure~\\ref{dm_temp} shows that this assumption fails in the case of cold fronts identified in our cluster simulations.  Our simulations show that dynamics of gas around the simulated cold front is much more complicated than that assumed in Vikhlinin et al.'s model.  In particular, the flow of gas around the front is not laminar and in general is not parallel to the front.  Vikhlinin et al., on the other hand, assumed a laminar flow similar to flow of gas around a blunt body.  In this case the gas flow lines are parallel to the front - the conditions ideal for development of the Kelvin-Helmholtz instabilities.  The velocity field shown in Fig.~\\ref{dm_temp} is very different and the conditions are probably not as conducive to the Kelvin-Helmholtz instability as those assumed in Vikhlinin et al.'s model.  On the other hand, the cold front in A3667 that they model has a much more regular appearance than the simulated cold front shown in Fig.~\\ref{chandra_map}.  But, such detailed comparison is difficult because processes discussed by Vikhlinin et al. would be poorly resolved in our simulations, which cannot resolve well perturbations smaller than 20~kpc.  It is also difficult to draw conclusions about how generic the differences are (i.e., whether the flow pattern assumed in Vikhlinin et al. can arise in simulated clusters).  These questions will be addressed in future studies using considerably higher resolution simulations and a large sample of simulated clusters, respectively.  To conclude, our results indicate that cold fronts are non-equilibrium transient phenomena arising in cluster mergers. They occur when the relatively cold gas of the merging subclump is moving fast with respect to the hotter gas of the main cluster.  Both simulated cold fronts discussed in the paper would be detectable in a {\\sl Chandra} observation of $\\sim 50$~ksec exposure, if the clusters were located at a moderately low redshift (e.g., $z\\lesssim 0.1$, see Fig.~\\ref{chandra_map}). Cluster mergers (especially minor) are common and the fact that we were able to identify several cold fronts in the simulations of only two clusters means that we can expect many more cold fronts to be discovered in the near future \\citep[e.g.,][]{sun02}. If the origin and properties of these features are understood, they can provide powerful constraints on the frequency of cluster mergers and properties of the intracluster medium. The analysis presented here is a first step in studying the origin and properties of cold fronts. The presented results should be useful in assessing the validity of the assumptions in the detailed modeling of cold fronts and developing more sophisticated models."
    },
    "0206/astro-ph0206375_arXiv.txt": {
        "abstract": "{ Pravdo et al (\\cite{PAH}) claimed that the phase resolved x-ray spectrum in Crab pulsar (PSR B0531+21) shows a spectral hardening at the leading edge of the first peak of its integrated profile (IP); this was a new and unexpected result. This article reanalyzes their data, as well as some other related data, and argues that the spectrum is as likely to be unvarying (i.e., neither hardening nor softening). ",
        "introduction": "Pravdo et al (\\cite{PAH}; henceforth PAH) studied the power law spectrum of Crab pulsar as a function of phase within its IP, in the energy range 5 KeV to 60 KeV, using $\\approx$ 86\\,000 periods of data obtained by the PCA detector aboard the RXTE x-ray observatory. Their photon arrival times had a resolution of 250 $\\mu$s, allowing them to obtain the spectrum in 60 on-pulse phase bins, each of duration $\\approx$ 335 $\\mu$s, out of a total of 100 phase bins within the IP. Their main conclusions are ``(1) the spectrum softens (i.e., the power law index increases) starting at the leading edge of the first peak until the intensity maximum of the first peak; (2) ``the spectrum hardens in the inter peak region; and (3) the spectrum softens throughout the second peak''. They point out that item (1) above is unexpected on the basis of existing theoretical models. They also point out that the power law index $\\alpha$ varying as the shape of first peak of the IP may be an important clue for understanding the details of the high energy emission of the Crab pulsar. Moreover, the increase in $\\alpha$ before the leading edge of the first peak of the IP coincides with the position of the radio precursor emission (Smith, \\cite{FGS86}; Lundgren et al. \\cite{LCM95}; Moffet \\& Hankins \\cite{MH1996}). All this adds to the significance of item (1) above.  This paper reanalyzes their data, as well as some other RXTE data that have better time resolution, and sufficient energy resolution within similar energy range. This paper argues that the spectral hardening is a weak effect in terms of the overall spectral variation across the IP; the spectrum is as likely to be unvarying in the leading edge of the first peak of the IP.  For each data file, obtained in the EVENT mode (XTE\\_SE), the Good Time Intervals (GTI) were obtained by using the MAKETIME tool on the corresponding XTE filter file; the selection criterion were (a) pointing offset less than 0.02$\\degr$, (b) elevation greater than 10$\\degr$, (c) all five PCUs to be switched on, and time since passage of RXTE satellite through the South Atlantic Anomaly (TIME\\_SINCE\\_SAA) to be greater than 30$\\degr$ or less than 0$\\degr$. Next the MGTIME tool was used to merge these GTIs with those in the second extension of each data file, with the AND option, to produce a final GTI file corresponding to each data file. These data/GTI file pairs were then input to the FASEBIN tool, along with the orbit file of the day, to obtain the phase resolved spectrum of Crab pulsar in 100 phase bins within the period. In the output of FASEBIN, the Crab nebula background was subtracted using an off-pulse phase range of 0.1, using the FBSUB tool.  \\begin{table}[h] \\begin{tabular}[t]{cccccccc} \\hline \\hline DATE & ID & FILE & EXP & CHN & $\\tau$ & ROFF & SHIFT \\\\ \\hline \\hline 1996 Apr 18 & 10202 & FS37\\_451ee10-451f318 & 193  & 10 & 30  & 0.02105 & -0.0007 \\\\ &       & FS37\\_451f6b0-451fc00 & 183  & 10 & 30  & 0.02105 & -0.0004 \\\\ 1996 Apr 20 &       & FS37\\_4536db0-4537a88 & 497  & 10 & 30  & 0.02105 & -0.0005 \\\\ 1996 May 2  & 10204 & FS37\\_463b300-463bf18 & 611  & 97 & 250 & 0.02110 & 0.0 \\\\ &       & FS37\\_463ca70-463d5a6 & 2014 & 97 & 250 & 0.02110 & 0.0 \\\\ &       & FS37\\_463e1e0-463eba8 & 464  & 97 & 250 & 0.02110 & 0.0 \\\\ 1996 Aug 23 & 10203 & FS3f\\_4f83ae0-4f84862 & 3310 & 10 & 4   & 0.02090 & -0.0003 \\\\ &       & FS3f\\_4f851d0-4f85ee1 & 3214 & 10 & 4   & 0.02090 & -0.0003 \\\\ &       & FS3f\\_4f86850-4f87561 & 3214 & 10 & 4   & 0.02091 & -0.0004 \\\\ &       & FS3f\\_4f87ed0-4f88be1 & 3214 & 10 & 4   & 0.02093 & -0.0004 \\\\ \\hline \\end{tabular} \\caption{ Columns 1,2 and 3 contain the date of observation, the ObsID of RXTE, and the name of the event mode data file. Column 4 contains the total duration of observation in seconds; this is obtained by summing the time intervals in the GTI file. In principle this is also supposed to be the time (divided by 100) found in the EXPOSURE column of the phase resolved spectrum; for unknown reasons this is not true for the ObsID 10204 data, but is true for the rest of the data. Column 5 contains the number of energy channels available within the energy range analyzed; this is 5.2 keV to 59.1 keV for ObsIDs 10202 and 10204, and 18.2 to 54.9 keV for ObsID 10203 (in Vivekanand \\cite{MV2001a} and \\cite{MV2001b}, which analyze data of ObsID 10203, the above energy range was wrongly stated to be 13.3 Kev to 58.4 KeV). Column 6 contains the time resolution of the data in micro seconds ($\\mu$s). Column 7 contains the offset time in seconds to be given to the FASEBIN tool. The last column contains the relative shift of the IPs of each data file, with respect to the IP of the data file FS37\\_463ca70-463d5a6, in units of pulsar phase; for comparison, the width of each phase bin is 0.01. } \\end{table}  Table 1 lists some information about the data analyzed in this paper. The data analyzed by PAH is FS37\\_463ca70-463d5a6, observed on 1996 May 2. They ignored the other two files observed on the same day, for ``dead-time and rate considerations'', although they state in their appendix that ``these 8 s data are usable for spectral analysis in principle, since the dead-time  is spectrally independent''. The rest of the data was chosen from the RXTE public data archive using the criterion: (1) the number of usable energy channels within 5 KeV to 60 Kev must be at least 10, for proper fitting of the power law spectrum, and (2) the time resolution should be much smaller than the width of each phase bin (which is $\\approx$ 335 $\\mu$s). PAH analyze data whose accuracy is comparable to the width of their phase bins; ideally the former should be much smaller than the latter. A fourth data file in ObsID 10202, and data in ObsID 40090, had the best time resolution (2 $\\mu$s), but $\\le$ 8 usable energy channels.  The output of the FASEBIN tool is the required phase resolved spectrum. This tool uses two parameters that can potentially affect further analysis -- the Crab pulsar period $P$ for the epoch of observation, and the time offset ROFF (column 7 of table 1) with respect to the radio ephemeris. The former is automatically obtained from the pulsar catalog, while the latter has to be inserted by hand. Now, the $P$ of epoch for Crab pulsar can be different from the catalog determined value owing to its glitching and timing activity; and an error of 3.5 nano second in the period used for folding can cause a total drift of 1 phase bin in one of the longer data files in table 1. Next, the data are obtained not in one contiguous data file, but in several files, each observed at different epochs. Therefore the correct time offset ROFF for each file is crucial to align in phase the data in those files; otherwise the $i^{\\mathrm{th}}$ (say) phase bin in each file might correspond to different true phase bins within the pulsar period.  To avoid these problems, two things were done. First, the ASCII version of the photon counts were obtained from the FITS version of each phase resolved spectrum. An IP was formed from this data, combining all energy channels. This was cross correlated with a standard IP (that of data file 10204:FS37\\_463ca70-463d5a6, which itself was first suitably shifted in phase for centering it). Any significant phase shift detected was added to ROFF, and the FASEBIN tool was run once again. In this iterative manner it was ensured that the residual phase shift (column 8 of table 1) is much less than the width of a phase bin (0.01).  Next, arrival times of photons in each data file above were converted to the solar system barycenter system (TDB) as described in Vivekanand (\\cite{MV2001a}, \\cite{MV2001b}), using the new FAXBARY tool (improved version in HEASOFT 5.1). Then the DECODEEVT tool was used to obtained the TDB times of each photon event, and its energy channel. From the ASCII version of this data, IPs were formed of the first and second halves of each data file above. These two were then cross correlated. Their relative phase shift was sufficiently small to ensure that the period used was sufficiently accurate for the current purpose (see Vivekanand \\cite{MV2001a}).  Penultimately, the response matrix for each data file was obtained using the PCARSP tool. Since the phase resolved spectrum output by FASEBIN is incompatible with PCARSP, a non-phase resolved spectrum was obtained for each data file using the tools SEFILTER and SEEXTRCT, using the corresponding GTI file.  Finally, the background subtracted phase resolved spectra were analyzed using the XSPEC tool. The spectrum was modeled as the power law $dN/dE = \\beta E^{-\\alpha}$ (photons per keV), where $\\beta$ is the normalization constant and $\\alpha$ is the power law index. Only energy channels lying between 2 keV and 60 keV were considered, since this is the RXTE/PCA energy calibrated range. The channels to ignore were obtained by looking at the corresponding response matrices using FDUMP. ",
        "conclusions": "This paper therefore concludes that the spectral hardening noticed by PAH, at the leading edge of the first peak of Crab pulsar's x-ray IP, is not a strong result, although PAH were probably not unjustified in arriving at such a conclusion from their own analysis. Other data, particularly that which was part of the data set of PAH but which they ignored, indicates that a constant $\\alpha$ is as likely a solution in contrast to a decreasing $\\alpha$.  Apparently the above spectral hardening has also been noticed by Massaro et al (\\cite{MCLT2000}) in the energy range 1.6 KeV to 300 KeV, using the BeppoSAX observatory. They claim in their discussion that regarding the first peak of the IP, they have `confirmed ...  that both the leading and trailing edges have harder spectra than the central bins'. In particular they note that in their figs. 3c and 3d, ``the central bins'' of the first peak ``have a softer spectrum than the wings''; these figures are similar to the bottom panel of fig.~2 here, but for energy ranges 10 keV to 34 keV, and 15 keV to 300 keV, respectively. However, a closer look at these two figures of Massaro et al (\\cite{MCLT2000}) indicates that this result may be be non-existent in their fig.~3d; $\\alpha$ appears to be more or less constant at phases below that of the first peak of the IP. Their fig.~3c certainly contains what might be interpreted as the spectral hardening noticed by PAH. If one were to take both these results seriously, one might be forced to the interpretation that the spectral hardening, if at all it exists in the BeppoSAX data, is likely to be a phenomenon confined to a narrow energy range (10 keV to 34 keV).  Figure 3b of Massaro et al (\\cite{MCLT2000}), energy range 1.6 keV to 10 keV, shows that $\\alpha$ may be constant at those phases where it decreases the most in their fig.~3c; and it shows a small bump where it is more or less constant in their fig.~3c. Finally their fig.~4a (15 KeV to 80 KeV range) appears to show a nearly unvarying $\\alpha$ at the leading edge of the first IP peak.  It therefore appears that the results of Massaro et al (\\cite{MCLT2000}) are also indicating the kind of confusion that has been presented in this paper, regarding the spectral hardening claimed by PAH.  At this stage one can certainly not rule out the spectral hardening claimed by PAH, particularly since another independent instrument claims to have noticed the same. However, one should also keep in mind the divergence of results in this regard, and also not rule out the possibility that $\\alpha$ may be more or less constant in the leading edge of the first peak of the IP. This is particularly important in view of (a) the lack of quantitative analysis in this regard, and (b) a recent theoretical study prefers a monotonically {\\bf increasing} $\\alpha$ at smaller phases within the IP, and neither a constant nor a decreasing $\\alpha$ (Zhang and Cheng \\cite{ZC2002}). To accommodate the spectral hardening of PAH, they need to invoke emission other than synchrotron radiation from secondary pairs.  Pravdo et al (\\cite{PAH}) claim that phase resolved x-ray spectrum in Crab pulsar ``softens throughout the second peak'' of the IP (their fig.~1). However fig.~\\ref{fig1} of this paper shows that this is also a weak effect, as do fig.~\\ref{fig4} and the bottom panel of fig.~\\ref{fig3}; in these figures the spectrum could be unvarying in the trailing edge of the second peak also. The above claim of PAH appears to be more evident in fig.~\\ref{fig2} of this paper. However, in the second panel of fig.~\\ref{fig3}, the spectrum appears to {\\bf harden} in the trailing edge of the second peak, opposite to the claim of PAH. Keeping in mind the lack of quantitative analysis, and the confusing visual results from different data sets, one can not rule out a constant $\\alpha$ in the trailing edge of the second peak of the x-ray IP of Crab pulsar."
    },
    "0206/astro-ph0206139_arXiv.txt": {
        "abstract": "{ We have used the optical lines of N II and Fe III to study the wind of the luminous blue variable \\object{P Cyg}. This was performed by applying a version of the Self Absorption Curve (SAC) method, involving few assumptions, to lines whose flux can be measured. A rather surprising result was obtained; the lines of more excited multiplets without blue shifted absorption components appear to be optically thick, while the lines of the most excited multiplets may show some indications of being optically thicker than the lines of less excited ones. Explanations of such effects are discussed, including possible inhomogeneities in the wind. ",
        "introduction": "The observational history of the hypergiant \\object{P Cygni} begun about 400 years ago, when it was discovered by Willem Blaeu as a nova-like object. Because of its spectral and photometric characteristics the star is considered to be the prototype of the P Cygni-type stars (PCT), as defined by Lamers (\\cite{L86}), and as one of the prototypes of the giant eruption Luminous Blue Variables (Humphreys \\cite{H99}). In addition, recent studies (Markova et al. \\cite{M01a}, de Groot et al. \\cite{GSG}) showed that over the last 15 years PCygni has behaved as a typical weak S Doradus variable in a short S Dor phase (van Genderen \\cite{vanGen01}). More information about this enigmatic object can be found elsewhere (Israelian and de Groot \\cite{IG99}, ASP Conf Ser. 233).   Stahl et al. (\\cite{S93}) reported the presence of a number of pure emission lines in addition to the dominant P Cygni-type lines in the optical spectrum of \\object{P Cygni}. Later Markova \\& de Groot (\\cite{MG97}) showed, comparing line-identification lists from various observational epochs, that the optical emission spectrum of \\object{P Cygni} was much richer and intense in the mid nineties of the nineteenth century, than 60 years ago. In particular, the authors noted that more than 70\\% of the pure emission lines appear to be, of recent origin. Among these are forbidden lines of Fe II, Ni II, Ti II Fe III and N II; high excitation lines of Fe III and  N II and lines of low and medium excitation of Si II.  If not due to observational selection caused by the continuous improvement of the signal-to-noise ratio of photographic spectra during the last century, this result might indicate the presence of a very long-term variation in the wind+photospheric properties of the star. In this context, it seems worthwhile to try to obtain additional information on the nature and the origin of P Cygni's emission spectrum  using methods not used up to now, such as for instance the SAC (Self-Absorption Curve) method developed by Friedjung and Muratorio (\\cite{FM87}). This method involves a semi-empirical analysis of emission line spectra of complex atoms and ions, without assuming detailed models for the objects emitting these spectra. By means of the SAC method valuable information (e.g. the self absorption effects, the level population laws and population anomalies) has been obtained for the Fe II emission spectrum of many stars of different types, such as  AG Car (Muratorio and Friedjung \\cite{MF88}), the VV Cep star KQ Pup (Muratorio et al \\cite{M92}), while procedures for using the method is described by Baratta et al (\\cite {B98}). Use of the SAC method can require however a few hypotheses about the relative populations of levels inside the same spectroscopic term and the distribution of self-absorption relative to that of the emission of different lines in different parts of the line emitting region. The second of these hypotheses was in any case not needed in the present work.   The main purpose of our study is to apply the SAC method to the Fe III and N II emission spectra of \\object{P Cygni}. In this way we can hope to gain additional information about the nature and the origin of these spectra and thus to get a deeper insight into the physics of the star itself. The observational material is described in Section 2, where a number of  problems concerning the derivation of the net emission equivalent widths of the studied lines are also discussed and resolved. Section 3 presents the results obtained through the SAC analysis while Section 4 deals with the interpretation of these results. ",
        "conclusions": "The origin of the emission spectrum of \\object{P Cygni} was partly discussed by Wolf and Stahl (\\cite{W85}) who suggested that the upper terms of multiplets 115 and 117 of Fe III were pumped by two lines of He I. Pumping by two UV transitions in He I ($\\lambda\\lambda$ 522, 537), followed by downward cascades, was suggested by Markova and de Groot (\\cite{MG97}) to explain the Fe III emission lines of multiplets 113, 114, 118 and 119. In addition, the authors noted that the N II multiplets with the highest upper term excitation potential ($\\sim$25 eV) might be populated by dielectronic recombination, but the temperature of the wind seems to be  too low for this mechanism to work. Stahl et al (\\cite{S93}) however suppose that the highest excitation N II multiplets are excited by an unknown pumping mechanism.  The optical thickness of the more excited Fe III and \\\\ N II lines without blueshifted absorption components, established by us, is quite surprising, while the indications of an increase of SAC slope with increasing excitation, if confirmed in future work, would be in fact counter-intuitive. At the same log$(gf \\lambda)$ more excited lines are expected to be in any case optically thinner. A conspirancy of differential non-LTE effects appears unlikely as an explanation, especially when the points for a multiplet do not show large deviations from a curve. However we should note, that it is hard to rigorously test for such an effect, as both the upper and lower levels of the observed optical Fe III and N II lines could in principle show deviations from LTE. In objects showing many Fe II emission lines, the lower levels of strong lines are very often in LTE and in that case one only needs to test for effects depending on the upper level.  Another explanation, which could be suggested, will not work either, if the increase of curve slope at the same log$(gf \\lambda)$ for more excited multiplets is confirmed in future work. One might suppose that the more excited lines are only emitted from certain parts of a spherically symmetric wind. In that case the absence of blueshifted absorption components for excited multiplets might be explained by a large source function, so the emission in the lines per unit surface area, would be larger than that of the photosphere. In that case extra emission instead of absorption would be present on the blue side of the more excited lines, resulting in a blueshift of the mean radial velocity, as was in fact observed by Markova and de Groot (\\cite{MG97}). Such a blueshift could be particularly important if part of the receding material was behind and so occulted by the photosphere, but the symmetry of the lines, noted by Markova and de Groot, would suggest that such occultation is small. In any case if the wind has spherical symmetry, the less excited lines could not be optically thinner, because those lines would also be emitted in the region of emission of the more excited lines, where they should not be optically thinner than those excited lines.  If we exclude the previous explanations, we may need to conclude that the wind of P Cyg is anisotropic. Optically thick excited lines with no blue shifted absorption components would in that case be mainly formed in clouds, which either have a large source function or which do not occult much of the photosphere. The optically thinner lower excitation lines with blueshifted absorption, could then be mainly formed in large regions of the wind, covering the photosphere, with only a small optically thicker contribution to their emission from the clouds. The physical reason for such a situation is however not immediately clear. Heating due to shock waves might occur in the clouds, producing extra ionizing radiation, followed by more recombinations and cascades plus pumping to extremely excited levels Such a situation might also be a way of explaining the presence of the most excited N II lines. Theoretical work is required to see whether this is a viable option.  We must also finally point out that effects such as an increase in self-absorption curve slope for more excited multiplets, would be hard to find in detailed spectral syntheses, based on spherically symmetric wind models. Semi-empirical methods can be quite powerful in searches for unknown physical processes, present in relatively complex situations."
    },
    "0206/astro-ph0206413_arXiv.txt": {
        "abstract": "We report the discovery of 39 Faint High Latitude Carbon Stars (FHLCs) from Sloan Digital Sky Survey commissioning data. The objects, each selected photometrically and verified spectroscopically, range over $16.6 < r^* < 20.0$, and show a diversity of temperatures as judged by both colors and NaD line strengths. Although a handful of these stars were previously known, these objects are in general too faint and too warm to be effectively identified in other modern surveys such as 2MASS, nor are their red/near-IR colors particularly distinctive.  The implied surface density of FHLCs in this magnitude range is uncertain at this preliminary stage of the Survey due to completeness corrections, but is clearly $>$$0.05$~deg$^{-2}$. At the completion of the Sloan Survey, there will be many hundred homogeneously selected and observed FHLCs in this sample.  We present proper motion measures for each object, indicating that the sample is a mixture of extremely distant ($>100$~kpc) halo giant stars, useful for constraining halo dynamics, plus members of the recently-recognized exotic class of very nearby dwarf carbon (dC) stars.  The broadband colors of the two populations are indistinguishable. Motions, and thus dC classification, are inferred for 40-50\\% of the sample, depending on the level of statistical significance invoked.  The new list of dC stars presented here, although selected from only a small fraction of the final SDSS, doubles the number of such objects found by all previous methods. The observed kinematics suggest that the dwarfs occupy distinct halo and disk populations.  The coolest FHLCs with detectable proper motions in our sample also display multiple CaH bands in their spectra. It may be that CaH is another long-sought low-resolution spectroscopic luminosity discriminant between dC's and distant faint giants, at least for the cooler stars. ",
        "introduction": "Although stars with prominent $C_2$ in their spectra have been observed for more than a century, faint high-latitude carbon stars (hereafter FHLCs), where here we arbitrarily define ``faint\" as $R>13$, prove to be of very current and special interest for a variety of oddly unrelated reasons.  Such objects are rare: certainly $<10^{-5}$ of random stellar images prove to be C stars. Thus FHLCs are, for example, rarer than QSOs at a given magnitude. They are also not particularly easy to discover: although the cool N-type stars (with apologies to \\citet{kee93} for the outdated nomenclature) do have very red colors, the considerably more numerous R and CH stars do not.  Although some FHLCs are found serendipitously, the majority of past discoveries have been due to objective prism surveys such as those at Case \\citep{san88}, Michigan \\citep{mac81}, Kiso \\citep{soy99}, Byurakan \\citep{gig01}, and Hamburg/ESO \\citep{chr01}. Recent attempts at automated photometric selection of FHLCs have met with some success for the very red, cool N stars \\citep{tot98, tot00, iba01}, but again the warmer, more numerous FHLCs have still proven difficult to select autonomously \\citep{gre94b}. Although $\\sim7000$ galactic carbon stars are known \\citep{alk01}, the sum of all the heterogeneous investigations discussed above has probably yielded only a few hundred FHLCs.  Why are we interested in finding these rare FHLCs, especially as essentially all are too faint for the high dispersion spectroscopic analysis that has been at the core of the study of red giants thus far? It has become clear in the past decade that the FHLC population consists of two totally distinct, physically unrelated classes of objects which (confusingly) share remarkably similar colors and (at least at moderate resolution) spectra. Both of these two classes are interesting.  Some fraction of the FHLCs are exactly what they appear to be: distant, luminous evolved giants in the halo.  There have been a small handful of previous hints that this population extends to rather astonishing distances. For example, \\citet{mar84} serendipitously found one such star at $d\\sim100$ kpc. Clearly the presence of a brief-lived phase of stellar evolution at these galactocentric distances poses interesting questions of origin:  could there be star formation in the distant halo, or in infalling gas? Are these the most luminous members of previously-disrupted dwarf satellites?  Moreover, aside from the question of the origin of the luminous FHLCs, they make splendid halo velocity tracers \\citep{mou85, bot91}, as at these huge distances they almost surely encompass the entire dark matter halo, and the very sharp $C_2$ band heads make radial velocity determinations straightforward even at modest sized telescopes.   The remainder of the FHLCs are perhaps even more exotic. They exhibit large proper motions \\citep{deu94}, and in some cases parallaxes \\citep{har98}, that place them at main sequence luminosity ($M_V\\sim10$).  These so-called ``dwarf carbon stars\", hereafter dC's, should be an oxymoron, as there should be no way for $C_2$ to reach the photosphere prior to the red giant phase. For 15 years, precisely one such star was known, G77-61 \\citep{dah77, dea86}, but \\citet{gre91} and \\citet{gre94a} showed that these are in fact a surprisingly common subclass of FHLCs, unnoticed in the past simply as most have $R>16$. A recent review of the dC stars has been given by \\citet{gre00}.  The dozen or so known previous to this work are all at $d<100$~pc, a volume that contains not a single giant C star.  Therefore, contrary to the conclusions of 100 years of classical astronomical spectroscopy, the overwhelming numerical majority of stars with $C_2$ in their spectra are in fact the previously unknown dwarfs, not giants! Current thinking is that the $C_2$ in dC's was deposited in a previous episode of mass-transfer from a now invisible, highly-evolved companion.  In this respect the dC's are probably similar to barium stars, and in particular to the so-called ``subgiant CH stars\" \\citep{bon74}. Those objects, which despite their names at least occasionally have near main sequence luminosity, are presumably slightly too warm to show strong $C_2$ despite the inference of $C/O>1$.  Aside from the usual invisibility of the evolved companion, great age for the dC stars is also implied by the extraordinarily metal poor composition inferred for the prototype, G77-61 \\citep{gas88}. Thus the ultimate significance of these stars might be to call attention to otherwise elusive Population~III objects (see also the discussion of \\citet{fuj00}). Furthermore, if an early generation of stars is responsible for reionization of the intergalactic medium, as now seems increasingly likely \\citep{mad99, fan01, bec01}, then the early Universe achieves a heavy element abundance already substantially above that inferred for the dC prototype, perhaps implying that objects such as these are actually pregalactic \\citep{ree98}.   Despite the totally different nature of the giant and dwarf C stars, the spectra and colors of the disparate classes are frustratingly similar. Indeed, although some preliminary photometric and spectroscopic luminosity criteria have been suggested \\citep{gre92, joy98}, dC's are currently identified with complete confidence only if they show detectable parallax or proper motion, thus ruling out membership in the distant halo.  The wide areal coverage ($10^4$~deg$^2$), faint limiting magnitude ($m\\sim22$), precision five-color photometry ($\\sim0.02$~mag), and highly multiplexed spectroscopic capabilities of the Sloan Digital Sky Survey \\citep{yor00} provide the opportunity to identify large numbers of new FHLCs, providing far larger samples than currently available, both for use as halo dynamic probes and to elucidate the nature of the engimatic dC's. In addition to merely extending the catalogs of both classes, major goals of this work are also to develop and/or refine photometric and spectroscopic luminosity discriminants, hopefully apparent at low to moderate resolution, and to understand the ratio of dwarfs to giants in a magnitude limited sample.  Here we report initial results of a search for FHLCs in the SDSS commissioning data. This subset of SDSS data is very similar, although not quite identical to, the SDSS Early Data Release (EDR), and the reader is referred to \\citet{sto02} for a detailed description of those data and their reduction. Our results therefore utilize only $\\sim5\\%$ of the eventual Survey, but should serve to illustrate the potential of SDSS to elucidate many issues related to the FHLC problem. This paper concentrates on methods of selection, and an overview of the first photometric, spectroscopic, and astrometric results, to give the reader an understanding of the nature of SDSS FHLC data. A more lengthy analysis of astrophysical implications will appear in later publications. ",
        "conclusions": "Despite their rarity and the relatively benign colors of the majority of the objects, large numbers of FHLCs can be efficiently selected by the SDSS. The current sample, although small compared with the ultimate end product of the Survey, already provides interesting information on a variety of FHLC issues. At the completion of the Survey, the homogeneously selected FHLC sample will for the first time be sufficiently large that statistics are no longer dominated simply by the size of the catalog, though a negligible fraction of SDSS resources are applied to this problem.  The SDSS has already fulfilled the previously undemonstrated expectations that a sufficiently sensitive and efficient selection technique for FHLCs would yield the (formerly) exotic dwarf carbon stars in copious numbers. For the first time we constrain the ratio of carbon dwarfs to giants in the $16 < r^* < 20$ range: it is at least near to unity, and possibly considerably larger. Preliminary kinematic analyses imply that there are distinct halo and disk dwarf populations.  Lacking the positive detection of proper motion, separation of the dwarfs from giants for FHLCs remains problematic. We find no single photometric or (low resolution) spectroscopic diagnostic that applies to all objects, but suggest the addition of one further weapon to the arsenal: in sufficiently cool FHLCs, the presence of strong CaH bands is an effective dwarf indicator.  Numerous unusual FHLCs have also been identified for further study."
    },
    "0206/astro-ph0206249_arXiv.txt": {
        "abstract": "{ We present high-resolution X-ray spectroscopic observations of M87 with the Reflection Grating Spectrometers on {\\it XMM-Newton}. We detect strong K-shell line emission from N, O, Ne, Mg, some emission from He-like Si, a fully resolved set of Fe L-shell emission spectra, and some emission from C.  The angular intensity distributions of the strong emission lines are detectably resolved on scales $(15-160) \\arcsec$.  The gas in the inner arcmin of M87 has a multi-phase structure, as indicated by the similarity of the emission line profiles of Fe L shell ions with widely separated ionization potentials. The global Fe L spectrum is approximately consistent with an isothermal plasma at $kT_e \\sim 1.8$ keV, in addition to a component with a temperature distribution appropriate to an isobaric cooling flow, but with a minimum temperature cutoff of $kT_{\\rm min} \\approx 600$ eV.  The behaviour of this cooling-flow component is qualitatively similar to what is seen in other cooling flow clusters. Finally, we do not find any strong evidence for a spatial variation in abundances due to resonance scattering redistribution in the inner arcminute of the core. ",
        "introduction": "The giant elliptical galaxy M87, its active nucleus, and its halo have been the subject of intensive studies at all wavelengths (see e.g., R\\\"oser \\& Meisenheimer 1999). The extended diffuse emission has been mapped in neutral atomic and molecular gas, and in the highly-ionised hot gas that is revealed through X-ray emission.  X-ray observations with the {\\it Einstein} Observatory first found evidence for a mass of rapidly cooling X-ray emitting gas at the core of the system. Estimates for the density and temperature of this gas suggested the presence of a `cooling flow' with a total mass deposition rate of $\\sim 10 M_{\\odot}$ yr$^{-1}$ (Fabricant et al. 1980; Stewart et al.  1984; Fabian et al. 1984). Direct spectroscopic evidence for the presence of radiatively cooling gas was obtained with the Focal Plane Crystal Spectrometer on {\\it Einstein} (Canizares et al. 1979) and the Solid State Spectrometer (Lea, et al. 1982).  The advent of high-sensitivity X-ray imaging and spectroscopy with the {\\it Chandra} and {\\it XMM-Newton} Observatories now provides the opportunity for a detailed study of the physical conditions in the cooling gas and its interaction with its environment. High spatial resolution images of M87 have been obtained with {\\it Chandra} (Wilson \\& Yang 2001).  Medium resolution {\\it XMM-Newton} spatially resolved spectroscopy has been published by B\\\"ohringer et al. (2001), Belsole et al. (2001), Molendi \\& Pizzolato (2001) and Matsushita et at (2001).  For objects of moderate angular extent (up to approximately 1~arcmin), the Reflection Grating Spectrometers (RGS) on {\\it XMM-Newton} have the unique capability to provide high resolution X-ray spectroscopy coupled with some spatial resolution in the cross-dispersion direction. This combination of capabilities is uniquely suited to conducting a sensitive study of the thermodynamic properties of the cooling gas. Here, we describe preliminary results of such a study, based on data obtained during the {\\it XMM-Newton} Performance Verification phase. ",
        "conclusions": "We have presented high resolution soft X-ray spectroscopy of M87. In brief, the \\xmm RGS spectrum shows fully spectrally resolved emission from C, N, O, Ne, Mg, and Fe. The angular structure of the intensity in the strong emission lines is resolved on scales of order \\gtae 15 $\\arcsec$, removing much of the confusion associated with the limited spectral resolution of CCD spectrometers. We do not find any evidence for resonance scattering in the central one arcminute.  The gas in the inner $\\sim 10$ kpc has a multiphase structure, as indicated by the similarity of the spatial profiles of the emission lines of Fe L ions with widely different ionization potentials. An analysis of the overall ionization balance indicates that gas at all temperatures down to $kT_{\\rm min} \\approx 0.6$ keV is present, with an emission measure distribution approximately equal to that predicted for isobaric radiative cooling.  Additionally, a strong approximately isothermal component of $kT_e \\approx 1.8$ keV is present, which gives rise to emission from higher Fe ionization states (e.g., Fe XXIV), whose distribution is much wider than that of the multiphased component, as seen in the CCD images (Molendi \\& Pizzolato 2001; Matsushita et al. 2001).  The global deposition rate of the rapidly cooling gas is about $\\Mdot \\approx 6 M_{\\odot}$ \\, yr$^{-1}$ in the inner arcmin.  In these respects, the core of M87 resembles that of other cooling flow clusters.  We do not find evidence for resonance scattering redistribution in the strong emission lines. Detailed calculations of the effect of radiative transfer in M87 (Mathews et al 2001) predict that the line of sight average optical depth of the O VIII line is $<$0.5, consistent with our results. However, the Fe XVII lines with higher oscillator strengths are predicted to be marginally optically thick in the absence of turbulence.  We do note that both our analysis and the reanalysis of the EPIC data have reduced the apparent O abundance deficit in the centre of M87, originally noted by B\\\"{o}hringer et al. (2001); this reduces the need for an inclusion of radiative transfer effects in the modelling of the M87 spectrum. If anything, the accurate determination of the metal abundance depends on the level of the continuum, which in turn depends on the contribution from the active nuclear/jet emission. More confident measurements will have to wait for data from which the non-thermal emission will be subtracted, instead of modelled.  Resonance absorption scattering has been recently found in a poor cluster of galaxies (NGC 4636; Xu et al. 2001). Apart from the apparent differences between the two sources (e.g., density and temperature), both of them host a radio galaxy. The radio galaxy in NGC~4636 (Birkinshaw \\& Davis 1985; Stanger \\& Warwick 1986) is much weaker than that in M87: NGC~4636 has a steep spectrum, with a core power at 2~cm of $6.6 \\times 10^{19} {\\rm W \\, Hz^{-1}}$ (Nagar et al. 2000). Both observational facts may imply that it is an old or a `dead' radio source. However, as is becoming clearer, thanks to recent numerical simulations (e.g., Br\\\"uggen et al. 2002), even `dead' radio galaxies can stir up, disrupt the intergalactic medium, and introduce turbulence. Therefore, a direct comparison between the two sources and the high resolution X-ray results, is not straightforward and more radio and X-ray observations with high angular resolution are required to address these issues. It is expected though that turbulence in the gas will reduce the net optical depth. Such turbulence could arise as the radio jets impinge onto the intergalactic medium. We defer the detailed discussion of the relative optical depths in the different lines and the effects of turbulence to a later paper which will model such effects directly.  Our analysis of the RGS data of M87 finds near-solar abundances of C and N, Ne and Mg, sub-solar for O, and 0.8 solar for Fe. We emphasise that the absolute values for the abundances have to be regarded with some caution, because they are based on fitting a parameterised model with limited flexibility.  As pointed out by many authors (e.g., Matsushita et al. 2000, Buote 2000) CCD spectra are not of sufficiently high resolution to obtain robust chemical abundances for complex thermal plasmas. In particular if there are sharp abundance gradients and multi-temperature gas the CCD data are degenerate, and several solutions can exist with different temperature and abundance values. The XMM CCD spectra for M87 (Finoguenov et al. 2002; Gastaldello \\& Molendi 2002) show both abundance gradients and complex thermal structure and therefore their interpretation is open to modelling uncertainties.  Both studies use rather different models for the temperature distribution, but arrive at similar values for the abundances except for oxygen. Gastaldello \\& Molendi (2002) point out that since the bulk of the emission in M87 is locally isothermal and shows a rather smooth temperature gradient, the abundances derived from the \\xmm CCD data should not be very sensitive to the assumed fitted temperature distribution; given the high signal to noise of the XMM CCD data, all acceptable solutions should produce essentially the same distribution of emission measure with temperature.  Our thermal modelling assumptions differ from those of both Gastaldello \\& Molendi (2002) and Finoguenov et al. (2002) but we do not believe that they should result in different abundances. The oxygen and Fe abundances we derive are the same as Finoguenov et al. (2002) and Gastaldello \\& Molendi (2002) averaged over the central 1~arcmin region.  However the abundances of Mg and Ne, which depend on detailed modelling the Fe L shell lines in the CCD data (but not in the RGS data) are rather different with the RGS results, both being ~(50-70)~percent higher. This comparison graphically illustrates the difficulty of accurately measuring the abundances of these elements in the temperature range for which the Fe L shell lines are strong.  The derivation of the Mg abundance is of particular importance: the M87 stellar spectra show a very high apparent Mg/Fe ratio of \\gtae 2:1 (Terlevich \\& Forbes 2002) and a super-solar Mg abundance which is completely incompatible with the RGS data of Mg/Fe $\\sim$1 and Fe less than solar.  Thus, very little of the observed X-ray emitting gas can originate from stellar mass loss, contrary to the expectations of the observed stellar density and normal stellar evolutionary theory. As originally pointed out by Lowenstein \\& Mathews (1991) a normal type I supernova rate combined with stellar mass loss would produce a 3-5 times solar Fe abundance in the gas. As discussed in detail by Awaki et al. (1994) the ASCA data did not detect such a high Fe abundance, raising the serious issue as to what has happened to the products of type I SN in the central regions of giant elliptical galaxies. These results enhance and deepen the mystery. The sub-solar O, Ne, Mg and Fe abundances are incompatible with the stellar data even if all the heavy elements in the gas originate from stellar mass loss.  We also know that M87 is not unique. The abundance pattern in NGC~4636 (Xu et al. 2001) is extremely similar and preliminary analysis of the RGS data for another giant elliptical NGC~533 (Peterson et al. 2002) shows similar results. We thus conclude that this issue is a serious one, and it needs further attention.  The \\xmm RGS results pose two problems regarding the M87 chemical abundances: i) why is the overall Fe abundance so low, and ii) why is the chemical composition apparently so different from the stars.  We speculate that the first problem can only be solved if either the true type I SN rate is considerably reduced, or the yield of Fe per unit supernova is also reduced. Otherwise, Fe may be lost from the central regions of M87 [by ram pressure stripping for example, (Stevens et al. 1999)], or the ISM gas is diluted due to the accretion of cluster gas, but both seem unlikely. The second issue can only be resolved if either the inferred stellar abundances are in error, or if the stellar spectra do not represent the true abundance distribution of the stars. If for example, there are two distinctive stellar populations in elliptical galaxies with different metal abundances, and the optical data register only the abundances of the metal rich group. Support for the chemical inhomogeneity of the stars in elliptical galaxies comes from deep VLT and HST observations, of NGC~5128 for example, (Harris \\& Harris 2000) which shows a wide range in metallicity of the stars."
    },
    "0206/astro-ph0206294_arXiv.txt": {
        "abstract": "A temporal analysis was performed on a sample of 100 bright short GRBs with T$_{\\rm 90}$$<$$2\\,$s from the BATSE Current Catalog along with a similar analysis on 319 long bright GRBs with T$_{\\rm 90}$$>$$2\\,$s from the same catalog. The short GRBs were denoised using a median filter and the long GRBs were denoised using a wavelet method. Both samples were subjected to an automated pulse selection algorithm to objectively determine the effects of neighbouring pulses. The rise times, fall times, FWHM, pulse amplitudes and areas were measured and their frequency distributions are presented. The time intervals between pulses were also measured. The frequency distributions of the pulse properties were found to be similar and consistent with lognormal distributions for both the short and long GRBs.  The time intervals between the pulses and the pulse amplitudes of neighbouring pulses were found to be correlated with each other.  The same emission mechanism can account for the two sub-classes of GRBs. ",
        "introduction": "It has been recognised that GRBs may occur in two sub-classes based on spectral hardness and duration with T$_{\\rm 90}$$>$$2\\,$s and T$_{\\rm 90}$$<$$2\\,$s \\citep{kmf:1993,nsb:2000,paciesas:2001}.  The bimodal distribution can be fit by two Gaussian distributions to the logarithmic durations \\citep{mhlm:1994}. A variety of statistical methods have been applied to the temporal properties of the long GRBs with T$_{\\rm 90}$$>$$2\\,$s. It is important to compare the temporal properties of the long and short GRBs to determine the similarities and differences between the two classes in an objective way. Detailed temporal analyses have been performed on a large sample of  short and long bright GRBs. The results from the long sample \\citep{quillig:2002} can be used as templates for comparison with a similar analysis of short GRBs. ",
        "conclusions": "Samples of short and long bright GRBs have been denoised and analysed by an automatic pulse selection algorithm. The results show that in both cases the distribution of the properties of isolated pulses and time intervals between all pulses are similar and compatible with lognormal distributions.  The same mechanism seems to be responsible for both long and short GRBs and may be attributable to the internal shock model."
    },
    "0206/astro-ph0206157_arXiv.txt": {
        "abstract": "We present results for $J$ and $K_s$ near-IR imaging data on a large sample of 88 galaxies drawn from the catalogue of Impey et al. (1996).  The galaxies span a wide range in optical and IR surface brightness and morphology (although they were drawn from a catalog constructed to identify low surface brightness galaxies, hereafter LSBGs).  They were also selected to include very low and high HI mass galaxies in order to ensure that they span a wide range of evolutionary states.  The near-IR data unveils many features of LSBGs not seen before in the optical.  First, a high fraction of the observed LSBGs are very luminous in the near-IR, indicating that they have a well developed old stellar population, and that older LSBGs are more frequent in the universe than data from optical bands suggested. Second, the near-IR morphologies are often quite different than seen in the optical.  Many diffuse LSBGs that are apparently bulgeless when observed in blue bands, instead exhibit nuclei in $J$ and $K_s$ bands.  Third, we find significant trends between the near-IR morphologies of the galaxies and their ratio of HI mass to near-IR luminosity.  Fourth, we find no trend in disk surface brightness with absolute magnitude, but significant correlations when the bulge surface brightness is used.  Finally, we find that the formation of a bulge requires a galaxy to have a total baryonic mass above $\\sim\\!10^{10}\\msun$.  A wide variety of other correlations are explored for the sample. We consider correlations among morphologies, surface brightnesses, near-IR colors, absolute magnitudes, and HI masses.  In addition, using previous results by \\citet{bell2001}, we convert the galaxies' near-IR luminosities to stellar masses, using color-dependent stellar mass-to-light ratios.  This allows us to consider correlations among more fundamental physical quantities such as the HI mass, the stellar mass, the total baryonic mass, the gas mass fraction, the mass surface density, and the metallicity (via the highly metal sensitive color index $J-K_s$).  We find that the strongest of our correlations are with the ratio of HI mass to total baryonic mass, M$_{HI}$/M$_{baryonic}$, which tracks the evolutionary state of the galaxies as they convert gas into stars, and which ranges from 0.05 up to nearly 1 for the galaxies in our sample.  We find strong systematic trends in how the metallicity sensitive $J-K_s$ color becomes redder with decreasing M$_{HI}$/M$_{baryonic}$, as would be expected for ``closed box'' models of chemical enrichment.  However, the increased scatter with increasing gas mass fraction and decreasing galaxy mass suggests that gas infall is increasingly significant in the gas rich lower mass systems.  We argue that the overall range in $J-K_s$ color argues for at least a factor of 20 change in the mean stellar metallicity across the mass range spanned by our sample.  We also see strong trends between M$_{HI}$/M$_{baryonic}$ and central surface density, suggesting that increased star formation efficiency with increasing gas surface density strongly drives the conversion of gas into stars. ",
        "introduction": "Low surface brightness galaxies (LSBGs) have been a subject of increasing interest in the last two decades.  Although an initial study by Freeman (1970) suggested that the central surface brightnesses of disk galaxies in the Hubble sequence (Sa-Sb-Sc) fell within a rather narrow range, \\citet{disney76} pointed out that the lack of low surface brightness galaxies may in fact be a selection effect due to the difficulty in discovering galaxies of very low surface brightness, as had previously been recognized by \\citet{zwicky57}.  Indeed, many surveys have since turned up large numbers of LSB galaxies.  In practice, ``low surface brightness'' has come to mean galaxies whose central surface brightness is fainter than 22.0 mag arcsec$^{-2}$ in the $B$ band-pass (i.e. more than $1\\sigma$ outside of the narrow range defined by \\citet{freeman1970} of $\\mu_0 = 21.7 \\pm 0.3 B$-mag/arcsec$^{2}$).  However, while many galaxies have been found below the ``Freeman value'' of surface brightness, disk galaxies with surface brightnesses that are significantly higher than those of normal spirals do not appear to exist. This is shown most clearly in \\citet{courteau96}, who find that there is a rather well-defined upper cut-off at a central surface brightness in $R$ of $\\sim 20.08$ mag arcsec$^{-2}$.  It was originally thought that all galaxies with a low surface brightness were early or late-type dwarfs \\citep{vandenbergh59}. However, radial velocity observations by \\citet{fisher75} showed that some galaxies with low surface brightness are actually quite large and luminous \\citep{mcgaugh92, bothun97}, for example LSBGs like Malin 1 type.  Indeed, most surveys have revealed LSBGs with disk scale lengths comparable to the Milky Way (although these angular diameter limited surveys are strongly biased toward finding the physically largest galaxies).  Based upon surveys to date, LSBGs do seem to have different properties than those of brighter galaxies (possibly indicating different evolutionary histories).  The observed trends suggest that LSBGs are in large part a continuous extension of the population of normal disk galaxies, reaching to higher angular momenta and lower masses \\citep{dalcanton97}.  Previous studies \\citep{longmore82} showed that LSBGs are much more gas rich (in terms of $M_{gas}/M_{stars}$) and bluer than ``normal'' late type galaxies \\citep{mcgaugh92, deblok95, sprayberry95}.  When combined with their observed low gaseous metallicity \\citep{mcgaugh94, deblok97}, the evidence suggest that LSBGs have had a historically low star formation rate and are relatively unevolved systems compared to the bulk of normal spirals. HI observations \\citep{vanderhulst93, deblok96} show that LSBGs have extended disks with low gas surface densities and high M/L ratios, further confirming the unevolved nature of LSBGs.  Even external influences seem to have had little effect in speeding up the evolution of LSBGs \\citep{mcgaugh97}.  Despite the tremendous progress in this field, there are still significant uncertainties in many of the above results.  First, while the data suggests a young stellar age for LSBGs overall, the spread in the integrated optical colors of LSBGs is large, suggesting that they may actually have traveled along diverse evolutionary paths.  In particular, because of the low surface brightness of the underlying population, only a small fraction of the total number of stars is needed to make the colors significantly blue, and thus a small change in the star formation rate can drastically change the apparent mean age of the stellar population as inferred from {\\emph{optical}} observations.  Second, while emission line measurements have revealed the metallicity of the current gas phase, they have not measured the metallicity of the underlying stellar population. Given the large gas reservoirs observed in LSBGs and the possibility of substantial late time gas infall, it is not clear that the current gas phase metallicity is closely related to the actual metallicity of the stars.  Third, the total stellar mass in these galaxies (and their contribution to dynamics) is rather uncertain, given the large uncertainties in the stellar mass-to-light ratio at optical wavelengths.  Finally, recent modeling of LSBGs predict that from the assumption that blue LSBGs are currently undergoing a period of enhanced star formation, there should exist a population of {\\em red}, non-bursting, quiescent LSBGs \\citep{gerritsen98, beijer99, bell2000, vandenhoek2000}. These galaxies should then also be metal-poor and gas-rich, and share many of the properties of the LSBGs observed in the optical.  In all of these areas, we can greatly improve the observational constraints by undertaking a systematic study of the properties of LSBGs and normal spirals in the near-IR.  Combined near-IR and (eventually) optical observations are a perfect probe to help to clarify the true stellar masses of LSBGs, to disentangle the roles played by the age and metallicity, and to understand their place in the continuum with normal high surface brightness (HSB) spirals.  We expect that contrary to initial impressions, LSBGs should be readily detectable in the near-IR.  Although much of the data suggests that the mean stellar population of LSBGs are probably young, we know from the very red integrated colors of some LSBGs (such as the giant Malin-type objects) that many LSBGs contain a significant {\\em old population of stars}, and were therefore not formed recently.  This old population should dominate the near-IR (which is a good tracer of red giant and lower mass stars) and thus should be detectable with NIR observations. We also should be able to easily identify LSBGs with bulges, as these are among the reddest members of the LSBG population \\citet{mcgaugh92}. There are also indications that the most massive LSBGs, such as the super-giant galaxies Malin 1 and Malin 2, are redder than the majority of LSBGs, with luminosities that are comparable to those of normal spirals.  Thus, near-IR observations should be particularly sensitive to the most massive end of the LSBG population.  While valuable, previous studies of the IR properties of galaxies in the low surface brightness regime have been limited due to the difficulty of the observations.  \\citet{knezek94} observed a preliminary sample of LSBGs in $J$, $H$ and $K$, but the sample was very small and biased toward massive galaxies.  \\citet{bergvall99} have observed 14 blue LSBGs in $J$, $H$ and $K$ and combined with optical photometry, arriving to the result that many properties of LSBGs observed in the optical are reproduced in the near-IR. In particular, they observe the same morphologies in $B$ and $J$, measuring weak optical/near-IR color gradients, a fact which they interpret as a low dust content.  Near-IR observations have also been used to probe the stellar population differences between LSBGs and normal spirals. \\citet{bell2000} undertook a combined optical and IR study of 26 LSBGs, including 5 from his earlier thesis work in \\citet{bell99}, and examples of both normal blue LSBGs and the rarer class of red LSBGs. They compare LSBG average ages and metallicities with their physical parameters and find that blue LSBGs are well described by models with low, roughly constant star formation rates, whereas red LSBGs are better described by a ``faded disc'' scenario.  Among the larger population, strong correlations are seen between an LSBG's star formation history and its $K$-band surface brightness, $K$-band absolute magnitude and gas fraction.  These correlations are consistent with a scenario in which the star formation history of an LSBG primarily correlates with its surface density and its metallicity correlates with both its mass and its surface density.  Another larger study involving IR data of spiral galaxies was made by \\citet{dejong96}, but included only a very small number of true LSBGs.   The lack of much previous work, and the availability of relatively large IR array detectors, both open the possibility to systematically observe a large sample of LSBGs and normal spirals in the near-IR, to derive fundamental parameters related to the origin of these galaxies and to their relation to normal galaxies.  Based upon the large expected rewards and the limited existing data, a systematic near-IR study of LSBGs is needed in order to complement existing optical data. We undertake such a study in this paper, which is organized as follows. In \\S\\ref{samplesec} we present the sample used in this work. \\S\\ref{obssec} briefly describes our observations, data reductions, photometric calibrations, galaxy photometry, and surface photometry. In \\S\\ref{resultsec} we analyze and discuss the most important results, including galaxy photometry, central surface brightnesses and scale lengths, the relationship between HI-mass and near-IR luminosities, and the relationship between stellar mass-luminosity ratios and baryonic mass.  Finally, we conclude in \\S\\ref{conclusionsec}. ",
        "conclusions": "\\label{conclusionsec}  The near-IR photometry of low surface brightness galaxies presented in this paper constitutes one of the largest database published until now.  Because of our original selection of galaxies from the \\citet{impey96} catalog, our sample includes not just a very large number of low surface brightness galaxies, but a number of normal spiral galaxies as well.  The total sample studied provides good indication that an old, red population of low surface brightness galaxies does exist.  Many of the LSBGs analyzed exhibit clear bulges and have high surface brightness disks in the near-IR, even though they were originally identified in a catalog of optical LSBGs\\footnote{These bulge dominated LSBGs also tend to be red, as found by \\citet{bell2000} for the optically selected red LSBGs of \\citet{oneil97}}.  Morphologically, the sample spans the full range of spiral Hubble types, from Sa to Im.  While there are broad correlations between the galaxies' optical and near-IR morphologies, we find several cases of galaxies which appear to be late-type in the optical, but have prominent bulges in the near-IR (e.g.\\ LSBGs \\# 100, 338, 384, and 446 in Figures \\ref{mosaics1} and \\ref{mosaics2}).  In addition to the morphologies, a quantitative analysis of the surface brightness profiles of the galaxies suggests that our sample contains galaxies with disk surface densities comparable to normal spirals, as well as a very high fraction of truly low surface brightness (and/or low surface density) disks.  On the bright end, we find that the central surface brightnesses of the disks in our sample show a well defined cutoff brighter than $\\mu_J(disk)\\sim 17.5$ mag arcsec$^{-2}$ (independent of the $J$-band absolute magnitude).  We argue that this cutoff is physical, and reflects the maximum surface density at which disks are stable.  On the LSBG end, our sample includes a very high fraction of galaxies that are indeed low surface brightness even in the near-IR, with some reaching roughly 3.5 magnitudes fainter in surface brightness than the bright cutoff.  The disks also have sizes that are systematically smaller for lower mass galaxies.  As part of our analysis, we have calculated the total {\\emph{baryonic}} masses (M$_{baryonic}=M_{gas}+M_{stars}$) of galaxies in our sample.  For M$_{gas}$, we have taken the HI mass from the \\citet{impey96} catalog, and assumed that other gaseous components are negligible.  For M$_{stars}$, we have used the color-dependent stellar mass-to-light ratios of \\citet{bell2001} to translate the observed near-IR luminosity into stellar mass.  The mass-to-light ratios also allow us to convert observed surface brightnesses into stellar mass surface densities.  The resulting masses and surface densities allow us to consider trends involving fundamental physical quantities (such as the ratio of the gas mass to the total baryonic mass).  The total baryonic mass itself is probably a good indicator of the total mass of the galaxies, assuming that baryon and dark matter are evenly mixed on large scales.  Among the masses of the various components, we see strong correlations, many of which follow trends previously identified in the optical (for example, see compilations in \\citet{schombert2001}). Galaxies with large stellar masses tend to also have high HI masses, higher central surface densities, but small gas mass fractions (i.e.\\ are HI-poor).  Likewise, galaxies with small stellar masses tend to have low HI masses but have a larger fraction of their baryonic mass in the gas phase (are HI-rich, as for the LSBGs studied by \\citet{mcgaugh97}).  We note that even galaxies which have formed large masses of stars can still retain large reservoirs of neutral gas, even after rich episodes of star formation.  We find several strong morphological correlations with the masses of the galaxies.  There are systematic trends towards more irregular, diffuse morphologies and lower surface densities with increasing gas mass fraction (or alternatively, M$_{HI}$/L$_{K_s}$).  We also find that the central surface brightness of the bulge component varies strongly with the near-IR luminosity and HI mass, as well as with the stellar and the total baryonic mass (such that the bulges have higher surface brightness for larger values of the listed quantities).  Moreover, we find that our sample exhibits a ``floor'' in the baryonic mass required for bulge formation, at $10^{10}\\msun$; we have no galaxies with lower baryonic masses that have formed high surface brightness bulges.  We have also explored how the near-IR color varies within our sample. In order to interpret the colors, we have used stellar population synthesis models to argue that the $J-K_s$ color is highly metal sensitive, such that redder colors imply higher metallicities.  We argue that based upon the range of observed colors our sample spans a factor of more than 20 in metallicity (but less than a factor of 100). We find that galaxies have redder $J-K_s$ colors with increasing HI mass, stellar mass, and baryonic mass, suggesting that more massive galaxies have in general have more metal enriched stellar populations.  We find very tightest correlations between color and the gas mass fraction M$_{HI}$/M$_{baryonic}$ in particular.  These correlations are significantly tighter than previous optically determined trends, even in the redder $V-I$ color (see compilation in Figure 6 of \\citet{schombert2001}).  In a closed box model for chemical evolution, there should be a one-to-one relationship between the metallicity and the gas mass fraction.  At small values of M$_{HI}$/M$_{baryonic}$, this is essentially what we see, in that the observed scatter in the metal-sensitive $J-K_s$ color is comparable to our observational uncertainties.  This suggests that in the gas poor, bulge dominated massive galaxies, chemical evolution has proceeded along the canonical closed box pathway (possibly confined to the bulge, which dominates the luminosity and thus the colors of the galaxy).  We find a different behavior among the gas rich, disk dominated, lower mass galaxies.  In these systems there is large scatter around the relationship between color and M$_{HI}$/M$_{baryonic}$.  This suggests that not only is the enrichment of these galaxies on-going (given their high gas mass fraction), but that it is not proceding as expected for a ``closed box'' model of enrichment.  Instead, it is likely that these gas rich systems are continuing to experience episodes of gas infall.  These episodes of gas accretion first increase the ratio of gas to stars, and then dilute the metallicity of subsequent generations of stars.  Episodic accretion would also explain the stochastic star formation suggested by analyses of the optical colors of LSBGs \\citep{oneil2000,bell2001} and theoretical calculations \\citep{gerritsen98,jimenez98}.  Temporary increases above a gas surface density threshold for star formation could dramatically increase the star formation rate \\citep{kennicutt98}, triggering bursts of star formation and an increased rate of gas consumption. This change in the luminosity-weighted mean stellar age would produce additional scatter in the color, scatter which will be more dramatic in the optical.  Within our sample, we see strong evidence that the star formation efficiency is indeed surface density dependent.  We find strong correlations between the gas mass fraction M$_{HI}$/M$_{baryonic}$ and the central stellar surface density (i.e.\\ M$_{stars}$/pc$^2$), suggesting that galaxies with the highest initial gas surface density have also had the highest efficiency in turning gas into stars.  This link between surface density and star formation efficiency also suggests that the young mean stellar age of LSBGs (as deduced from luminosity-weighted optical colors) results principally from the galaxies being slow to convert gas into stars, rather than from their being slow to assemble (see also \\citet{bell2000}, \\citet{schombert2001}, \\citet{deblok98}).  However, the evidence for on-going gas accretion in lower luminosity galaxies suggests that this second effect does play some role as well.  For future work, we intend to extend the current work with the addition of optical colors from our on-going database of $B$ and $R$ photometry for the same galaxies presented here, as well as from other bands provided by The Sloan Digital Sky Survey. We shall compute dynamical masses, and current gas phase metallicities using also future spectroscopic data. This will allow us to perform more critical and accurate tests of the ideas presented in this work."
    },
    "0206/astro-ph0206011_arXiv.txt": {
        "abstract": "We study the hydrodynamical evolution of massive accretion disks around black holes, formed when a neutron star is disrupted by a black hole in a binary system. The initial conditions are taken from three--dimensional calculations of coalescing binaries. By assuming azimuthal symmetry we are able to follow the time dependence of the disk structure for 0.2~seconds in cylindrical coordinates $(r,z)$. We use an ideal gas equation of state, and assume that all the dissipated energy is radiated away. The disks evolve due to viscous stresses, modeled with an $\\alpha$-law. We study the disk structure, and in particular the strong meridional circulations that are established and persist throughout our calculations. These consist of strong outflows along the equatorial plane that reverse direction close to the surface of the disk and converge on the accretor. In the context of gamma ray bursts (GRBs), we estimate the energy released from the system in neutrinos and through magnetic--dominated mechanisms, and find it can be as high as $E_{\\nu}\\approx 10^{52}$~erg and $E_{BZ}\\approx 10^{51}$~erg respectively, during an estimated accretion timescale of 0.1--0.2 seconds. $\\nu \\overline{\\nu}$ annihilation is likely to produce bursts from only a short, impulsive energy input $L_{\\nu \\overline{\\nu}} \\propto t^{-5/2}$ and so would be unable to account for a large fraction of bursts which show complicated light curves.  On the other hand, a gas mass $\\approx 0.1-0.25 M_\\odot$ survives in the orbiting debris, which enables strong magnetic fields $\\approx 10^{16}$~G to be anchored in the dense matter long enough to power short duration GRBs. We highlight the effects that the initial disk and black holes masses, viscosity and binary mass ratio have on the evolution of the disk structure. Finally, we investigate the continuous energy injection that arises as the black hole slowly swallows the rest of the disk and discuss its consequences on the GRB afterglow emission. ",
        "introduction": "Accretion onto black holes has been considered as an efficient way to transform gravitational energy into radiation \\citep{s64,z64}, and is often thought to occur in the form of a disk, due to the angular momentum of the accreting matter. This almost certainly is the case in a variety of astrophysical systems, ranging from AGNs with very massive black holes to stellar mass binaries, analogous to the X--ray binaries known to contain neutron stars \\citep{king95}.  The flows in these disks may exhibit very different morphologies depending on the physical conditions present in the system. Accretion is generally thought to proceed through the transport of angular momentum from the inner to the outer regions of the disk, although the mechanism by which this is accomplished is not entirely clear. The parametrization introduced by \\cite{ss73} has allowed much progress to be made, without specifying the physics behind the viscosity responsible for angular momentum transport. Magnetohydrodynamical studies (analytical and numerical) appear to indicate that magnetic fields and their associated instabilities in disks can effectively generate a viscosity that would drive their evolution \\citep{bh91,hb91,h00,sp01,hk01,hk02}, with equivalent values for the $\\alpha$ parameter in the range 0.01--0.1 \\citep{bh98}.  Additionally, it has become clear that hydrodynamic processes can play an important role in the structure and evolution of disks, and that multi--dimensional, time--dependent computations are necessary to fully understand these effects, particularly since the flows can be quite complicated, often exhibiting strong variability, and composed of combinations of inflows/outflows of varying intensity \\citep{igu00}.  In this context we are motivated to study accretion disks around black holes hydrodynamically, particularly in what concerns central engines for cosmological gamma ray bursts (GRBs). The energetics of GRBs ($10^{50}-10^{52}$~erg are typically released in a few seconds) and the variability shown in their lightcurves (down to millisecond timescales) argues in favor of a compact source that produces a relativistic outflow, usually referred to as a fireball \\citep{rm92}. The complicated light curves can then be understood in terms of internal shocks in the outflow itself, caused by velocity variations in the expanding plasma \\citep{rm94,sp97,frw99,rf00,np02}. The durations range from $10^{-3}$~s to about $10^{3}$~s, with a bimodal distribution of short ($\\simeq 0.2$~s) and long bursts ($\\simeq 40$~s) \\citep{kouveliotou93,n96}.  One possible scenario, at least for the class of short bursts\\footnote{We note, furthermore, that those with detected afterglows are all in the long category (although see \\cite{lrg01} and \\cite{c02}).}, involves the coalescence of compact binaries containing a black hole (BH) and a neutron star (NS), or two neutron stars \\citep{ls76, bp86, bp91, eich89, nara92, moch93, katz96, klee98, pop99, rj99, sr02}. Such systems do exist, like PSR1913+16 \\citep{ht75}, and they will coalesce due to angular momentum losses to gravitational radiation, provided that the orbital separation is small enough. The binary coalescence rates are in rough agreement with the observed GRB rate \\citep{kalo01}. The typical dynamical timescale in such binaries immediately prior to coalescence ($\\approx$~ms) is much shorter than the observed burst duration, and so it requires that the central engine evolves into a configuration that is stable, while retaining a sufficient amount of energy to power the burst.  The formation of a BH with a debris torus around it is a common ingredient of both these scenarios, whose accretion can provide the release of sufficient gravitational energy $\\approx 10^{54}$ erg to power a GRB \\citep{rees99}. A fireball arises from the large compressional heating and dissipation associated with this accretion, which can provide the driving stress necessary for relativistic expansion \\citep{piran99,m01}. Possible forms of this outflow are kinetic energy of relativistic particles generated by $\\nu\\overline{\\nu}$ annihilation or an electromagnetic Poynting flux.  In either mechanism, the duration of the burst is determined by the viscous timescale of the accreting gas\\footnote{In the collapsar scenario, the burst duration is given by the fall-back time of the gas \\citep{w93,mw99}}, which is significantly longer than the dynamical timescale, thus accounting naturally for the large difference between the durations of bursts and their fast variability. Any instability (of hydrodynamical or magnetic origin) would presumably be reflected in the relativistic outflow. Strong magnetic fields anchored in the dense matter surrounding the BH would produce large amplitude variations in the energy release \\citep{u92, m97}. A weaker field would extract inadequate power; on the other hand a neutron torus, with its huge amount of differential rotation, is a natural site for the onset of a dynamo process that winds up the magnetic field to the required intensity \\citep{u94, kl98}. An acceptable model requires that the orbiting debris not be dispersed completely on too short a timescale, lasting at least as long as the characteristic duration of the burst.  The tidal disruption of a neutron star by a black hole would redistribute matter and angular momentum very rapidly. The key issue is then how long a sufficient amount of this matter survives to power a burst.\\\\  In this paper, we study the evolution of realistic disks resulting from dynamical coalescence calculations (see below \\S \\ref{method}), on timescales that are comparable to the durations of short GRBs (i.e. a few tenths of a second).  Recently, the steady state structure of similar disks has been examined by \\citet{nara01} and \\citet{km02}. Since we wish to investigate the evolution of the disk and its stability, no steady state assumptions are made regarding its structure. This allows us to compare the strength of the energy release by $\\nu\\overline{\\nu}$ annihilation relative to MHD coupling and to investigate their potential as a viable source for GRB production.  The hydrodynamics of black hole--neutron star coalescence has been considered in detail by \\citet{lk99a,lk99b} and \\citet{l00,l01}, using an ideal gas equation of state and varying its stiffness through the adiabatic index $\\Gamma$.  The simulations explored the effect of different initial binary mass ratios $q_{b}=M_{\\rm NS}/M_{\\rm BH}$ and the spin configuration of the neutron star, with respect to the orbital motion. The viscosity inside neutron stars is far too small to permit synchronization during the inspiral phase \\citep{koch92, bild92}, so it is reasonable to assume that the neutron star spin is negligible with respect to the orbital angular momentum. The exact equation of state at supra--nuclear densities is uncertain, but it would appear that the radius is largely independent of the mass \\citep{prak01}, so that in the polytropic approximation, one would assume $\\Gamma=2$. Coalescence calculations with this set of parameters were computed by \\citet[herafter L01]{l01}. In particular, we will use here the results of runs C50 and C31 in that paper, with $q_{b}=0.5$ and $q_{b}=0.31$ respectively (both runs used an index $\\Gamma=2$). ",
        "conclusions": "We have computed the dynamical evolution of massive accretion disks around stellar-mass black holes in two dimensions (with azimuthal symmetry), formed through the tidal disruption of a neutron star by a black hole in a close binary ($M_{\\rm BH}\\approx 4M_{\\sun}$ and $M_{disk}\\approx 0.3M_{\\sun}$). Our initial conditions are taken from the final state of three dimensional hydrodynamical calculations of the coalescence process, by averaging in the azimuthal direction. We use Newtonian physics, an ideal gas equation of state, and solve the equations of viscous hydrodynamics assuming an $\\alpha$ law for the viscosity coefficient. All the energy dissipated by the physical viscosity is radiated away (in neutrinos). The time evolution is followed for 0.2~seconds.  We find that after an initial transient of numerical origin, stemming from the fact that the 3D torus does not exhibit strict azimuthal symmetry due to the highly dynamical merging process (see Figures 2 and 7 in L01), the disk settles to a qualitatively steady state. Meridional circulations are promptly established, whose structure depends mainly on the value of the $\\alpha$ parameter. Most strikingly, there is an important motion of fluid from the inner regions of the disks to large radii, along the equatorial plane. The flow is directed towards the accreting black hole along the surface of the disk and in the equatorial region at small radii. The disks remain thick ($H/R \\simeq 0.5$) throughout the dynamical evolution, due to their large internal energy, with accretion rates on the order of one solar mass per second. The maximum densities decrease during our calculations, as there is no external agent feeding the disks, but remain at $\\simeq 10^{12}$~g~cm$^{-3}$, with corresponding internal energy densities $\\simeq {\\rm few} \\times 10^{30}$~erg~cm$^{-3}$.  We stress that the evolution of accretion disks such as these should be studied with time--dependent models, since the system is clearly not in a steady state, even from its inception. The circulation pattern seen in Figure \\ref{vEG} would certainly lengthen their lifetimes by moving matter to larger radii continuously, an effect that would otherwise be omitted. To illustrate how important this can be, we considered the total radial mass flux in the disk, composed of two parts at any given value of the radial coordinate $r$, $\\dot{M}(r)=\\dot{M}_{in}(r)+\\dot{M}_{out}(r)$ where \\begin{equation} \\dot{M}_{in}=2 \\pi r \\int_{v_{r}<0} \\rho v_{r} dz, \\; \\; \\dot{M}_{out}=2 \\pi r \\int_{v_{r}>0} \\rho v_{r} dz, \\end{equation} restricted to regions in which $v_{r}<0$ for $\\dot{M}_{in}$ and $v_{r}>0$ for $\\dot{M}_{out}$ respectively. By definition, $\\dot{M}_{in}<0$ and $\\dot{M}_{out}>0$. If there were no circulation in the disk and all matter moved radially inward one would have $\\dot{M}_{out}=0$ and $\\dot{M}=\\dot{M}_{in}<0$. A measure of how important the circulations are, and how much they would lengthen the lifetime of the disks can be obtained by calculating the fraction of the gas flowing radially in the disk that is actually moving toward the accretor, i.e. $|\\dot{M}_{in}|/(|\\dot{M}_{in}|+|\\dot{M}_{out}|)$. In the inner regions of the disks this ratio tends to unity, as can be seen from Figure~\\ref{vEG}. It decreases rapidly at larger radii, reaching about 1/3 for run E and 1/10 for run G at $r \\simeq 100$~km midway through the simulations (at $t=0.1$~s). Thus most of the radial flow of gas actually cancels out in the circulations, with a residual amount left over moving toward the black hole.  We now turn to the implications our calculations might have on models for the production of cosmological gamma ray bursts from coalescing compact binaries. The two main forms of energy release from the disk we consider are i) neutrino emission and ii) MHD flow, either through the Blandford--Znajek effect \\citep{bz77} or by means of a magnetized wind.  In the first case, we make a rough estimate for $L_{\\nu}$ and $E_{\\nu}$ from the dissipated energy because of viscosity (see Table~\\ref{table:evol}), as mentioned above.  We have made an estimate only for the total neutrino luminosity $L_{\\nu}$, and not for the {\\em annihilation} luminosity $L_{\\nu \\overline{\\nu}}$, which would determine if a relativistic fireball could be launched or not. The calculation of $L_{\\nu \\overline{\\nu}}$ requires the use of a more realistic equation of state, which we will explore in future work. For the time being, we note that, regardless of the efficiency of energy conversion from neutrino luminosity to annihilation luminosity --- which could be quite low, on the order of 1 per cent or less \\citep{rj99,pop99}--- it appears that the time--dependence of $L_{\\nu}$, which follows that of $\\dot{M}_{\\rm BH}$ (see Figure \\ref{mdotbz}a and Table \\ref{table:evol}), is such that neutrinos could only be responsible for a very short, almost impulsive energy release ($L_{\\nu}\\propto t^{-5/4}$, so $L_{\\nu \\overline{\\nu}} \\propto L_{\\nu}^{2}\\propto t^{-5/2}$), and thus would be unable to power a burst lasting several tenths of a second or more. Of course this does not mean that it would have a negligible impact on the structure of the burst itself. We note furthermore that in the detailed 3D calculations done by \\cite{rj99} for the evolution of thick disks following the coalescence of two neutron stars, the ``neutrinosphere'' is quite close to an isodensity surface at $\\rho=10^{11}$~g~cm$^{-3}$, which is lower than the maximum densities present in the disks computed here, even at late times (see Figures~\\ref{rhoEG} and \\ref{LS}). \\cite{sr02} have also performed 3D calculations of binary neutron star coalescence taking into account neutrino emission and scattering processes, finding as well that in the regions of highest density the material is opaque, in their case mainly because of scattering off heavy nuclei. Thus it is clear that a complete picture must include an appropriate formulation of neutrino transport.  For the magnetic--dominated case we must make some assumptions, since our simulations do not incorporate the effects of magnetic fields explicitly. For the field to be able to extract the binding energy of the torus, it should be anchored to it, and we assume its magnitude is directly related to the internal energy in the gas. For this purpose we show in Table~\\ref{table:evol} the internal energy density $\\rho c_{s}^{2}$ at $t=0.1$~s in the inner regions of the disk (at $r=1.25 r_{Sch}$ ($\\approx 15$~km for runs A and B, and $\\approx 20$~km for runs C through G) in the equatorial plane. It is of order $10^{30}$~erg~cm$^{-3}$, and even larger for the case with low viscosity ($\\alpha=0.01$). From this we compute an estimate for the Blandford--Znajek luminosity as \\begin{equation} L_{BZ}\\approx10^{50} a^{2} \\left(\\frac{M_{\\rm BH}}{3~M_{\\sun}}\\right)^{2} \\left(\\frac{B}{10^{15}~{\\rm G}}\\right)^{2}~{\\rm erg}~{\\rm s}^{-1} \\end{equation} where $a \\simeq 0.3$ is the Kerr parameter of the black hole and the magnitude of the magnetic field is computed using $B^{2}/8\\pi=\\rho c_{s}^{2}$ (this gives $B\\approx 10^{16}$~G in all cases). This is clearly the most optimistic scenario concerning energy release, in that it assumes that the magnetic field strength is at the equipartition value. The time evolution of $L_{BZ}$ is shown in Figure~\\ref{mdotbz}b for several runs.  For a larger viscosity, the gas in the disk drains into the black hole on a shorter timescale, and thus the drop in $L_{BZ}$ is much faster than for a low value of $\\alpha$ (in run G, $L_{BZ}\\approx 5 \\times 10^{51}$~erg~s$^{-1}$ is practically constant).  In principle, energy can be extracted over many dynamical timescales if magnetic fields can tap the rotational energy of the accretion disk and of the black hole. Here $L_{BZ} \\propto (t/t_0)^{q^{\\prime}}$ is the intrinsic luminosity of the central engine measured in the fixed frame, where $-5/4<q^{\\prime}<0$ (see Figure \\ref{mdotbz}b). For such a continuously fed fireball, a forward shock propagating into the external medium and a reverse shock moving back into the relativistic outflow will coexist on either side of the contact discontinuity. The latter may persist as long as a significant level of energy injection is maintained. The differential energy conservation relation for the self-similar blast wave can be written as $dE/dt =L_{BZ}(t/t_0)^{q^{\\prime}} -\\kappa^{\\prime}(E/t)$ \\citep{cps98}, where the first term denotes steady energy injection, and the second accounts for possible radiative energy losses of the fireball, with $q^{\\prime}$ and $\\kappa^{\\prime}$ being dimensionless constants. In the adiabatic case (no radiative losses) $\\kappa^{\\prime}=0$ and the Lorentz factor evolves as $\\Gamma^2 \\propto t^{-3}$ \\citep{bm76}.  For $q^{\\prime}\\neq -1-\\kappa^{\\prime}$, an analytical solution can be found: $E=[L_{BZ}/(q^{\\prime}+\\kappa^{\\prime} +1)](t/t_0)^{q^{\\prime}} + E_{\\rm imp}(t/t_0)^{-\\kappa^{\\prime}}$ for $t>t_0$, where $E_{\\rm imp}=E_{\\nu \\overline{\\nu}}$ describes the impulsive energy input. Here $t_0$ is the characteristic timescale for the formation of a self-similar solution, which is roughly equal to the time for the external shock to start to decelerate. For $t>t_0$, the bulk Lorentz factor of the fireball scales with time as $\\Gamma^2 \\propto t^{-m}$, with $m$ and $\\kappa^{\\prime}$ connected by $\\kappa^{\\prime}=m-3$, and $m=3$ for an adiabatic blast wave expanding in a constant density medium. In the observer frame, the arrival time at the detector $T$ is related to that in the fixed (laboratory) frame $t$ by $dT=(1-\\beta)dt$, and $T=\\int_0^{t} (2\\Gamma^2)^{-1}dt \\approx t/[2(m +1)\\Gamma^2]$ \\citep{fcrs99}.  The differential energy conservation relation in the observer frame is now given by $dE/dT ={\\cal{L}}_{BZ}(T/T_0)^{q} -\\kappa(E/t)$, and the integrated relation is \\begin{equation} E={{\\cal{L}}_{BZ} \\over \\kappa+q+1} \\left({T \\over T_0}\\right)^q T + E_{\\nu \\overline{\\nu}} \\left({T \\over T_0}\\right)^{-\\kappa},\\;\\;\\; T > T_0. \\end{equation} Here ${\\cal{L}}_{BZ}=2\\Gamma^2(t_0)L_{BZ}$, and $q=(q^{\\prime}-m)/(m+1)$, $\\kappa=\\kappa^{\\prime}/(m+1)$.  Since $\\kappa + q +1 = (q^{\\prime}+\\kappa^{\\prime}+1)/(m+1)$, the comparisons between $q$ and $-\\kappa -1$ in the following discussion are equivalent to the comparisons between $q^{\\prime}$ and $-\\kappa^{\\prime}-1$. At different times, the total energy of the blast wave may be dominated either by a continuous injection term or by the initial impulsive term \\citep[and references therein]{zm01}. Which of the two is dominant at a particular observation time $T$ depends both on the values of the two indices ($q$ and $-\\kappa -1$), and on the relative values of ${\\cal{L}}_{BZ}$ and $E_{\\nu \\overline{\\nu}}$.  If $q<-1-\\kappa$, the impulsive term always dominates since the first term is negative.  For an adiabatic blast wave expanding in a constant density medium (i.e. $m=3$ and $\\kappa=\\kappa^{\\prime}=0$), this condition yields $q^{\\prime}< -1$. If $q>-1-\\kappa$, the continuous term will eventually dominate (i.e. $E_{\\nu \\overline{\\nu}} \\ll {\\cal{L}}_{BZ}T_{0}$) over the second term after a critical time \\begin{equation} T_{c}=T_{0} \\times {\\rm max}\\left[1,\\left((\\kappa + q +1){E_{\\nu \\overline{\\nu}} \\over {\\cal{L}}_{BZ}T_{0}}\\right)^{1/(\\kappa+q+1)}\\right]. \\end{equation} This is the case considered in many pulsar central engines \\citep{u92, u94, dl98}. We then have $T_{c} \\approx T_{0}$ and the fireball is completely analogous to the impulsive regime with $E_{\\rm total} \\approx {\\cal{L}}_{BZ}T_{0}$. \\\\  If initially $E_{\\nu \\overline{\\nu}} \\gg {\\cal{L}}_{BZ}T_0$, the critical $T_{c}$ after which the continuous injection becomes dominant could be much longer than $T_{0}$, exerting a noticeable influence on the GRB afterglow. We expect this to be the case either when the magnetic energy is below equipartition ($\\rho c_{s}^2 > B^{2}/8\\pi$) or when the timescale for the formation of a self--similar solution is small ($T_{0} \\ll T_{c}$). Because of the strong pinch that develops, a narrow jet that delivers its thrust in a narrow solid angle, $\\Omega_{BZ}$, may be a common ingredient of strong rotating magnetic fields (see \\citet{mku00} for a recent review). We thus generally expect the magnetic outflow to be much more collimated than that produced by $\\nu \\overline{\\nu}$ annihilation ($\\Omega_{BZ} \\ll \\Omega_{\\nu \\overline{\\nu}}$) and the magnetic luminosity to be the dominant contribution. Even in this case, the impulsive term $E_{\\nu \\overline{\\nu}}$ may be responsible for either creating a cavity before the magnetized wind expands or for precursor emission. The detection of, or strong upper limits on, such features would provide constraints on the burst progenitor and on magnetar--like central engine models.  An external shock can occur at much larger radii and over a much longer timescale than in usual afterglows, if the environment has a very low density. This may be the case for GRBs arising from compact binary mergers that are ejected from the host galaxy into an external medium that is much less dense than the ISM assumed for usual models (where the particle density is $n \\sim 0.1-1$~cm$^{-3}$). It is commonly assumed that compact mergers occur outside the host galaxies because of the long inspiral, due to the emission of gravitational waves. We note that recent calculations of evolutionary tracks of binary systems containing massive stars show that the merger events can occur on much shorter timescales and thus still within the host galaxy, because the initial binary separation is small \\citep{bbk02}.  Our simulations would indicate that the central engine survives the initial, violent event that created it, and that it possesses enough energy to account for the energetics and durations of short GRBs. However, they clearly cannot tackle directly other relevant issues, mainly related to the evolution of the magnetic field, and its influence on the dynamics. Magnetic instabilities could make the disk lifetime much shorter by effectively increasing the viscosity. The amplification of the magnetic field may be self--limiting due to magnetic stress, which would cause disk flaring. The properties of the expected variability depend strongly on the details of the configuration of the disk corona generated by the magnetic field, which is removed from the disk by flux buoyancy \\citep{nara92, tp96, kl98}."
    },
    "0206/astro-ph0206337_arXiv.txt": {
        "abstract": "{We map a 50$'$ x 30$'$ area in and around the M17 molecular complex with the French submillimeter balloon-borne telescope PRONAOS, in order to better understand the thermal emission of cosmic dust and the structure of the interstellar medium. The PRONAOS-SPM instrument has an angular resolution of about 3$'$, corresponding to a size of 2 pc at the distance of this complex, and a high sensitivity up to 0.8 MJy/sr. The observations are made in four wide submillimeter bands corresponding to effective wavelengths of 200 $\\mic$, 260 $\\mic$, 360 $\\mic$ and 580 $\\mic$. Using an improved map-making method for PRONAOS data, we map the M17 complex and faint condensations near the dense warm core. We derive maps of both the dust temperature and the spectral index, which vary over a wide range, from about 10 K to 100 K for the temperature and from about 1 to 2.5 for the spectral index. We show that these parameters are anticorrelated, the cold areas (10-20 K) having a spectral index around 2, whereas the warm areas have a spectral index between 1 and 1.5. We discuss possible causes of this effect, and we propose an explanation involving intrinsic variations of the grain properties. Indeed, to match the observed spectra with two dust components having a spectral index equal to 2 leads to very large and unlikely amounts of cold dust. We also give estimates of the column densities and masses of the studied clumps. Three cold clumps (14-17 K) could be gravitationally unstable. } ",
        "introduction": "}  The submillimeter domain is particularly suited to characterizing dust properties in the interstellar medium. Dust emission in this spectral range is mainly due to big grains at thermal equilibrium (see, e.g., \\cite{desert90}), whose emission is usually modelled by the modified blackbody law, {\\it i.e.} by the temperature and the spectral index of the dust. The temperature of a molecular cloud is a key parameter which controls (with others) the structure and evolution of the clumps, and therefore, star formation. Thus spectral imaging of molecular clouds can provide useful information about their structure and evolution, especially if the dust emission parameters can be properly derived on top of submillimeter intensities. Mapping of star-forming molecular clouds, as well as other dusty regions, has been performed by the PRONAOS balloon-borne experiment (PROgramme NAtional d'Observations Submillim\\'etriques, see \\cite{serra01} or \\cite{ristorcelli98}). We present in this article the maps and analysis made from PRONAOS observations of the M17 star-forming complex.  The Messier 17 Nebula (also called Omega, Horseshoe or Swan Nebula) is an ionized region associated with a giant star-forming molecular cloud located at about 2200 parsecs from us (\\cite{chini80}) in the constellation of Sagittarius. This nebula has the largest known ionization rate in the Galaxy for star forming regions (see for example \\cite{glushkov98}). This large ionization rate is mainly due to the excitation of gas by young O-type stars. The M17 molecular cloud has been mapped in carbon monoxide emission by Lada (1976), showing the two most intense condensations in this cloud, usually called M17 North (N) and M17 Southwest (SW). This cloud is part of a giant molecular complex extending 170 pc to the southwest, along the Sagittarius spiral arm (CO emission shown by \\cite{elmegreen79}). The interaction of this giant molecular cloud with the H II regions is particularly visible in the M17 region, where an expanding shock front interacts with the gas clouds, and is thought to have fragmented the original molecular cloud (\\cite{rainey87}). The M17 SW cloud is the best studied region of the area, especially for the photon-dominated region (PDR) near the boundary of the H II region (to the northeast). The first far infrared observations of M17 were made by Low \\& Aumann (1970) and Harper \\& Low (1971). M17 SW was mapped in the mid and far-infrared by Harper \\etal (1976) and Gatley \\etal (1979). Wilson \\etal (1979) mapped the whole M17 cloud ({\\it i.e.} including M17 North) at 69 $\\mic$. The IRAS satellite (http://www.ipac.caltech.edu/ipac/iras/iras.html) provided infrared maps of the whole sky, which showed the distribution of the dust in the M17 complex, however without giving much information on the faintest and coldest regions of the complex. The recent JCMT measurement of Wilson \\etal (1999) provided precise CO maps of the M17 molecular cloud, as well as the measurement of Sekimoto \\etal (1999) on a larger area. In this context, submillimeter mapping can provide useful information about the dust properties, especially in cold regions, and give independent estimations of the ISM masses in this region. For these reasons, we have observed a large region including the M17 cloud and fainter areas, with the multi-band photometric instrument (SPM) of the PRONAOS balloon-borne experiment. The maps are 50$'$ by 30$'$ (about 30 pc x 20 pc at the distance of M17) at angular resolutions between 2$'$ and 3.5$'$. We present in Section 2 how the observations were made and the way we processed the data to make maps. Section 3 presents the resulting maps, and Section 4 presents a detailed analysis of these maps. ",
        "conclusions": "We observed a large (50$'$ x 30$'$, $\\approx$ 30 pc x 20 pc) area in and around the M17 molecular complex. We showed extended faint clumps (A, B, C, D) in the outskirts of the active star-forming area. Our study shows a large distribution of temperatures and spectral indices: the temperature varies roughly from 10 K to 100 K, and the spectral index from 1 to 2.5. The statistical analysis of the temperature and spectral index spatial distribution shows the anticorrelation between these two parameters. Indeed, we observe cold dust (10-20 K) with high indices (around 2), and warm dust ($>$ 20 K) with low indices (around 1 - 1.5), but we could not significantly find any cold place with low indices nor warm place with high indices. The investigations which we made to match the observed spectra, with two dust components having a standard spectral index, imply very large and unlikely amounts of cold dust, therefore we rather support a fundamental explanation for this effect. Laboratory measurements showed this anticorrelation in the submillimeter domain on grains for temperatures down to 25 K, but we would need laboratory results on temperatures down to 10 K in the submillimeter range to fully understand our observations. This anticorrelation effect was also shown in Orion (\\cite{dupac01}), and in other PRONAOS observations that are still being analyzed.  We estimated the column densities and masses of the observed regions by simply modelling the thermal emission of the grains from D\\'esert \\etal (1990) or Ossenkopf \\& Henning (1994). We derived a total mass of the M17 cloud of 31000 \\msol. Of course, the error bars are large and difficult to estimate properly, but there is a clear trend for three cold clouds (A, B, D) to be gravitationally unstable, and therefore to be pre-stellar candidates.  These observations could be sustained by sensitive continuum observations in the millimeter domain, especially for the faint clouds studied to the west of the M17 complex, in order to better constrain the spectral index measurement of the cold dust. Also, infrared observations at higher resolution could be useful to study the structure of these cold clouds."
    },
    "0206/astro-ph0206101_arXiv.txt": {
        "abstract": "Spectacular results of pulsar observations in X-ray and gamma-ray domains gave a new boost to theoretical models of pulsar magnetospheric activity. A challenging aspect of these efforts is that lightcurves and broadband energy spectra of the brightest HE sources exhibit unexpected richness of features, making each object almost a unique case to be interpreted with a custom-made model. This review offers our subjective account of how the models of high-energy radiation used within the framework of SCLF polar-cap scenarios tackle with these challenges. We describe major characteristics of all radiative processes relevant for modeling the observed features. Then we address successes as well as noticeable disadvantages of these models upon their confrontation with the available data. ",
        "introduction": "A large fraction of galactic X-ray and gamma-ray sources is associated with neutron stars, in particular with rotation powered pulsars (RPP). Spectacular results of observational campaigns of RPP with ROSAT, ASCA and CGRO, and recently with RXTE and CHANDRA induced a new wave of interest in theoretical models of pulsar magnetospheric activity. A challenging aspect of these efforts is that lightcurves and broadband energy spectra of the brightest HE sources exhibit unexpected richness of features making thus each object a unique case pending special treatment.  The pair creation paradigm is a pivotal element in virtually all models of magnetospheric activity of radiopulsars (which at the moment constitute a vast majority of known RPPs). Electron-positron pairs ($e^\\pm$-pairs) are thought to be responsible for radio emission observed in radiopulsars which is interpreted as the coherent curvature radiation of $e^\\pm$ plasma. Pairs can be produced in magnetospheric environments  either via photon absorption in a dense field of soft photons (photon-photon collision) or via photon absorption in a strong magnetic field. In either case a supply of hard gamma photons is required in order to fulfill stringent threshold conditions for pair creation. Not all of these photons would be subject to absorption; many will escape the magnetosphere without any attenuation. This leads us to expect that all radiopulsars (including low-$B$ millisecond objects) should be the sources of gamma radiation. It is up to theoretical models to show whether the expected flux of the radiation is interestingly high with respect to the sensitivity of recent and future gamma-ray telescopes. One should keep in mind, however, that the pair creation paradigm remains just a paradigm for the time being: recently it was questioned by \\cite{jessner}, who presented arguments that for a wide range of surface temperature, magnetic field strength and spin period, the electrons supplied by neutron star surface via thermionic and field emission will screen out the accelerating electric field ${\\cal E}_\\parallel$, limiting it to a residual value. In consequence, no favourable conditions would exist for magnetic pair production.  This review offers our subjective account of how the models of high-energy radiation do perform within the framework of polar-cap models with space-charge-limited-flow (SCLF). First, we introduce all ingredients of polar-cap models with SCLF and we describe major characteristics of all radiative processes relevant for physical conditions in pulsar magnetospheres. Then we address successes as well as noticeable disadvantages of these models upon their confrontation with the data. ",
        "conclusions": "High-energy astrophysics of pulsars was challenged by unexpected richness of spectral and temporal properties found for the brightest gamma-ray pulsars. Numerous modifications (both, minor and major) to the existing models of magnetospheric activity are being invented to accommodate at least some of these properties. It is clear, however, that we actually need  good quality high-energy data for much weaker sources. Only then it will be possible to asses on statistical grounds the significance of those properties (the argument used by many authors, recently reiterated strongly by \\cite{b2001}). Chances are that it may have happen in a few years time.  The planned observatory GLAST \\citep{glast} will be superior to  EGRET  in two aspects. First, its sensitivity at 10~GeV will be more than two orders of magnitude better than that of EGRET. Second, it will reach energy of 300~GeV, closing thus for the first time the energy gap between satelite and ground-based  observatories (HE and VHE, respectively). The MAGIC Telescope \\citep{blanch} -- a 17 m diameter Imaging Air Cherenkov Telescope  (IACT) -- is expected to operate with sensitivity about three orders of magnitude higher at 10~GeV than EGRET. Its advanced technology will make possible to cover energy range  between 10~GeV and 1~TeV, and to reach $\\sim 50\\,$TeV in the Large Zenith Angle mode. Energy ranges of GLAST and MAGIC will overlap over more than one decade in energy.  We expect that the question of whether millisecond pulsars emit gamma-rays at all, will be answered relatively easily with these high-sensitivity observatories, verifying at first the status of  J0218$+$4232.\\\\ In the context of numerous positive detections  anticipated at hard gamma-rays, inclusion of viewing geometry effects will be of particular importance to properly interpret the observed shapes of spectral turnovers at high-energy ends.  To tackle high-quality data of the future in the most effective way, numerical 3D codes should be developed, capable of fast tracing the development and propagation of electromagnetic cascades in variety of non-dipolar realizations of magnetic field structure. Relaxing the centered-dipole assumption as a first step  would be very much in line with recent conclusions from the soft X-ray data analysis for neutron stars (Pavlov et~al.~\\cite{pavlov2002}). Optical to gamma-ray properties of Vela (section~\\ref{problems}) are in our opinion a signature of a strongly non-axisymmetric hollow cone of radiation around the magnetic axis. Strong deviations of the actual magnetic field structure from a pure dipole at the stellar surface may well be responsible for inducing this axial asymmetry already at the site of electron acceleration. (Recently, multipolar character of open field lines has been considered in a quantitative way by Gil et~al.~\\cite{gil} to justify radiopulsars with vacuum gap solutions in pulsars with superstrong magnetic fields.) An advantage of incorporating  multipolar components to the polar-cap models is twofold: \\\\ 1) magnetic field strength may locally be lowered slightly, opening thus windows for GeV photons to escape without magnetic attenuation; postulating high-altitude accelerators would be then unnecessary.\\\\ 2) directional characteristics of high-energy radiation would change notably with respect to the dipolar axis at higher altitudes, where radio emission is generated."
    },
    "0206/astro-ph0206047_arXiv.txt": {
        "abstract": "We have carried out near-infrared spectroscopic observations of 23 very low-luminosity young stellar object (YSO) candidates and 5 their companions in Heiles Cloud 2, one of the densest parts of the Taurus molecular cloud. Twelve objects were confirmed as YSOs by Br$\\gamma$ feature. The effective temperatures of the YSOs and of the companions are estimated from the 2.26 $\\micron$ feature, the 2.21 $\\micron$ feature, and the H$_{2}$O band strengths. Detailed comparisons of our photometric and spectroscopic observations with evolutionary tracks on the HR diagram suggest some objects to be very low-mass YSOs. ",
        "introduction": "Recent optical and near-infrared photometric studies have revealed faint populations of young stellar object (YSO) candidates in low-mass star forming regions (\\cite{Comeron}; \\cite{Strom95}; \\cite{ITG}, hereafter ITG; \\cite{Barsony}; \\cite{Oasa99}), in intermediate-mass star forming regions (\\cite{Aspin}), and even in high-mass star forming regions (\\cite{Kaifu}; \\cite{Lucas00}; \\cite{Oasa02}). Their faintness is suggestive of low-mass; they may be very low-mass young stars near the stellar/substellar boundary, young brown dwarfs, or even free floating planets.  From the photometric observations alone, however, it is impossible to simultaneously determine the mass and age of a YSO. Near-infrared spectroscopy of YSO candidates is necessary to overcome this difficulty (\\cite{Greene95}; \\cite{Luhman97}; \\cite{LuhmanRieke}; \\cite{Luhman98}; \\cite{Wilking}; \\cite{Cushing}; \\cite{Lucas01}). Spectra of some faint YSO candidates exhibit the absorption features distinctive to late spectral type, implying young brown dwarfs. While such very low-mass objects may be ubiquitous in star-forming regions, detail of the formation process of such objects, for example mass function, object density, and disk property, is still unknown.  \\citet{Briceno98} have carried out an optical search for very low-mass YSOs in the L1495, L1529, L1551, and B209 regions in the Taurus molecular cloud. From photometry with spectroscopy, they found 9 new YSOs in the clouds, half of them have very late spectral types, implying very low-mass objects (0.05 \\MO -- 0.25 \\MO). \\citet{Luhman00} further investigated very low-mass YSOs in the same clouds by combining optical imaging with near-infrared spectroscopy. They found that the mass function of the YSOs in these regions has a peak at 0.8 \\MO and is relatively flat between 0.1 \\MO and 0.8 \\MO range. In these papers, however, because the targets for the spectroscopy were fully or partially selected based on the optical color-magnitude diagram, the sample may not be complete especially for embedded very low-mass objects.  ITG conducted a near-infrared survey of the central \\timeform{1D}$ \\times$ \\timeform{1D} region of Heiles Cloud 2 in the Taurus molecular cloud, one of the best-studied low-mass star forming regions, with a limiting magnitude of 13.4 mag in the $K$-band. Fifty YSO candidates were identified by their intrinsic red color on the ($J-H$, $H-K$) color-color diagram, following the scheme discussed by \\citet{Strom93}. Successive high-resolution imaging survey discovered 5 companion candidates around the YSO candidates \\citep{ITN}. Faintness of some YSOs and their companion may be an indicative of the low-mass of the objects.  We describe here near-infrared spectroscopic follow-up of these faint YSO candidates in the Heiles Cloud 2. The observations are described in \\S 2, and data reduction procedures in \\S 3. In \\S 4, we derive effective temperatures of the YSOs mainly from the 2.21 $\\micron$ feature, the 2.26 $\\micron$ feature, and the H$_{2}$O absorption band. We also calculate photospheric luminosity of the YSOs from the previous photometry, then plot the YSOs on the HR diagram. ",
        "conclusions": "\\begin{enumerate}  \\item We have carried out near-infrared spectroscopic observations of 23 very low-luminosity YSO candidates and 5 of their companions in Heiles Cloud 2 in the Taurus molecular cloud. Near-infrared spectroscopy is essential to characterize objects in the color-color diagram. Out of the 28 objects, 5 objects have Br$\\gamma$ in emission and 7 have \"flat\" spectra over the Br$\\gamma$ feature. We conclude that these 12 objects are indeed YSOs.  \\item Compiling near-infrared spectra of dwarfs and giants taken by us as well as those in the literature, the ratio of the 2.26 $\\micron$ feature to the 2.21 $\\micron$ feature turns out to be a good indicator of the effective temperature for M type stars.   \\item The effective temperatures of the YSOs are determined. These objects are cool (T$_{\\mathrm{eff}} <$ 4000 K) YSOs.  \\item The mass of these YSOs is estimated from the HR diagram with recent evolutionary tracks. Some objects are very low-mass (0.1 \\MO -- 0.3 \\MO) YSOs.  \\item The age of these YSOs appears to be 10$^{5}$ -- 10$^{7}$ years. However, the deduced age depends on evolutionary track models.   \\end{enumerate}  ~\\\\  ~\\\\ ~\\\\  We are grateful to T. Geballe, T. Kerr, and Y. Oasa for help with the UKIRT observations, and H. Terada, N. Kobayashi, and B. Potter for the Subaru observations. We thank T. Tsuji, and T. Nakajima for discussions on spectral features of late type stars. We also thank our referee, B. Wilking, for many helpful comments. Y. I. is supported from the Sumitomo Foundation. A part of this study was supported by the UK-Japan collaboration fund from the JSPS. The United Kingdom Infrared Telescope is operated by the Joint Astronomy Centre on behalf of the U.K. Particle Physics and Astronomy Research Council. Subaru Telescope is operated by the National Astronomical Observatory of Japan.   \\appendix"
    },
    "0206/hep-th0206193_arXiv.txt": {
        "abstract": "It is well known that modifications to the Friedmann equation on a warped brane in an anti de Sitter bulk  do not provide any low energy distinguishing feature from standard cosmology. However, addition of a brane curvature scalar in the action  produces effects which can serve as a distinctive feature of brane world scenarios and can be tested with observations. The fitting of such a model with supernovae Ia data (including SN 1997ff at $z\\approx1.7$) comes out very well and predicts an accelerating universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206498_arXiv.txt": {
        "abstract": "A strong correlation is reported between gamma-ray burst (GRB) pulse lags and afterglow jet-break times for the set of bursts (seven) with known redshifts, luminosities, pulse lags, and jet-break times.  This may be a valuable clue toward understanding the connection between the burst and afterglow phases of these events.  The relation is roughly linear (i.e.\\ doubling the pulse lag in turn doubles the jet break time) and thus implies a simple relationship between these quantities. We suggest that this correlation is due to variation among bursts of emitter Doppler factor.  Specifically, an increased speed or decreased angle of velocity, with respect to the observed line-of-site, of burst ejecta will result in shorter perceived pulse lags in GRBs as well as quicker evolution of the external shock of the afterglow to the time when the jet becomes obvious, i.e. the jet-break time.  Thus this observed variation among GRBs may result from a perspective effect due to different observer angles of a morphologically homogeneous populations of GRBs.  Also, a conjecture is made that peak luminosities not only vary inversely with burst timescale, but also are directly proportional to the spectral break energy.  If true, this could provide important information for explaining the source of this break. ",
        "introduction": "Only recently, with the discovery of afterglows and in turn, redshifts for a handful of gamma-ray bursts, has there been progress in trend spotting within the seemingly chaotic variety of gamma-ray burst shapes and sizes.  \\citet{nmb00} discovered an anti-correlation between the isotropic peak gamma-ray luminosity, $L_{pk}$, of GRBs and the pulse lag, $\\Delta t$.  This lag is the time delay of the arrival of a burst pulse in the BATSE detector low energy channels compared to its arrival in the high energy channels.  Similarly \\citet{frr00} and also \\citet{rlfr+00} have shown that a measure of the variability of GRB lightcurves correlates with this peak luminosity.  Most recently \\citet{fksd01} have shown that the isotropic gamma-ray energy, $E_{iso}$, is anti-correlated with the jet-break time, $\\tau_j$.  The jet-break time is when the afterglow lightcurve changes (typically seen as a break) its decay rate, which is thought to be a manifestation of the finite opening angle of the jet.  As demonstrated in \\citet{sg01} these correlations are closely related and are likely manifestations of the same physical effect.  As discussed in the next section, we find an unexpectedly tight relationship between spectral lags and jet-break times.  Thus we argue that transitivity suggests that $L_{pk}$, $E_{iso}$, $\\Delta t$ and $\\tau_j$ are all interrelated by power-laws.  In \\citet{jay00,jay01} it was argued that the lag-luminosity relationship, $L_{pk}$ vs.\\ $\\Delta t$, derives from kinematics: the variation in velocity of the relativistic ejecta with respect to the observer.  In particular, the Doppler factor, dependent upon the speed and angle of the emitter with respect to the observer, will increase observed luminosity and decrease observed timescales.  In \\citet{sg01} we argue that all of these relationships originate from kinematic variations among bursts. ",
        "conclusions": ""
    },
    "0206/astro-ph0206451_arXiv.txt": {
        "abstract": "Approximately four thousand light curves of red variable stars in the Large Magellanic Cloud (LMC) were selected from the 2.3-year duration MOA database by a period analysis using the Phase Dispersion Minimization method. Their optical features (amplitudes, periodicities, position in CMD) were investigated. Stars with large amplitudes and high periodicities were distributed on the only one strip amongst multiple structure on the LMC period-luminosity relation. In the CMD, the five strips were located in the order of the period. The stars with characterized light curves were also discussed. ",
        "introduction": "The multiple and complicated period-luminosity relation for red variables in the LMC had been discovered using the microlensing database (Wood, Alcock, Allsman et al. 1999, Wood 2000). Although the Mira sequence (Feast, Glass, Whitelock, and Catchpole 1989, Hughes \\& Wood 1990) have been remarked as a distance indicator, such multiplicity is fatal for use as a distance indicator because their characteristics of each strip have not been revealed actually. The MOA (Abe, Allen, Banks et al. 1997; Hearnshaw, Bond, Rattenbury et al. 2000) database of the LMC obtained by large-scale photometry is quite appropriate to study the above problem. Not only related with the multiplicity of the period-luminosity relation, the photometric properties must be studied carefully to reveal the nature of AGB variables. Some interesting results of the study of the MOA database is presented. ",
        "conclusions": ""
    },
    "0206/astro-ph0206445_arXiv.txt": {
        "abstract": "We present a time-dependent photoionization code that combines self-consistently metal evolution and dust destruction under an intense X-ray UV radiation field.  Firstly, we extend the mathematical formulation of the time-dependent evolution of dust grains under an intense radiation flux with the inclusion of the process of ion field emission (IFE). We determine the relative importance of IFE with respect to X-ray and UV sublimation as a function of grain size, intensity and hardness of the incident spectrum. We then combine the processes of dust destruction with a photoionization code that follows the evolution of the ionization states of the metals and the relative radiative transitions. Our code treats, self-consistently, the gradual recycling of metals into gas as dust is sublimated away; it allows for any initial dust grain distribution and follows its evolution in space and time.  In this first paper, we use our code to study the time-dependent behaviour of the X-ray and optical opacities in the nearby environment of a Gamma-ray Burst, and show how the time variability of the low-energy and high-energy opacities can yield powerful clues on the characteristics of the medium in which the bursts occur. ",
        "introduction": "The realization that Gamma-ray Bursts (GRBs) are of cosmological origin has made them rank among the most energetic astrophysical phenomena known to us.  Even though the total energy output is much smaller than that of other bright sources on the sky, such as quasars, their luminosity can be much higher.  As a result, while the surrounding region affected by their radiation is much smaller than the corresponding region for QSOs, the effects of the interaction can be observed on a very short time-scale, comparable with the duration of the burst and its longer wavelength emission.  The X-ray and UV radiation accompanying a GRB alters the equilibrium of the medium in its close vicinity by heating and photoionizing it, and vaporizing dust grains. The time-variability of absorption lines and photoionization edges in GRB spectra has been discussed by Perna \\& Loeb (1998), B\\\"ottcher et al.  (1999) and Lazzati, Perna \\& Ghisellini (2001).  The destruction of dust by the intense UV flash produced by the reverse shock accompanying GRBs has been treated by Waxman \\& Draine (2000) and Draine \\& Hao (2002; DH in the following), while the effect of X-rays on dust particles has been discussed by Fruchter et al. 2001).  DH (see also Draine 2000) also computed in detail the absorption spectrum due to the vibrationally excited levels of the H$_2$ molecule.  The time-dependent effects resulting from the interaction of a GRB and its longer wavelength radiation with the environment can be a powerful diagnostics of the type of environment in which the bursts occur. Lazzati \\& Perna (2002) showed how the time-dependent X-ray extinction is sensitive to the density profile in the close environment of the bursts, which is different in the various progenitor scenarios for GRBs. On the other hand, as discussed before, GRBs, even though have a much shorter lifetime than QSOs, they can be far more luminous.  As such, they can allow to probe regions that are not accessible to QSO absorption studies due to the high absorption. It has in fact been observed (Fall \\& Pei 1993) how studies of Damped Lyman $\\alpha$ (DLA) absorbers done through QSO absorption spectra are biased against observing the densest systems. And it might very well be the case that the outer regions probed through these studies do not constitute a fair sample of what the properties of these high-redshift systems are. Therefore, absorption studies with GRBs can constitute a wonderful complement to absorption studies made with QSOs, and give us a more complete picture of what the properties of high redshift galaxies really are. Along these lines, it should be mentioned a study of metal column densities and dust content made with three GRB spectra by Savaglio, Fall \\& Fiore (2002). By comparing the inferred column densities and dust depletion amounts with those derived in DLAs from QSO absorption studies, they inferred that the GRBs were probing denser and dustier regions than the ones probed by QSOs.  A proper evaluation of the densities and dust content of the GRB close environment requires a knowledge of how GRBs affect their environment, modifying its properties, and in particular how they influence gas and dust in their surrounding.  So far, the evolution of metals and dust under the influence of an intense radiation field has always been treated separately. In this paper, we present a time-dependent photoionization code that incorporates metals and dust evolution in a self-consistent way. Firstly, we improve the photoionization code developed by Perna, Raymond \\& Loeb (2000) and Perna \\& Raymond (2000) with the addition of a proper treatment of radiative transfer in optically thick media and the inclusion of the opacity of the H$_2$ molecule.  We then extend the mathematical formulation of the time-dependent evolution of dust grains under an intense radiation flux with the inclusion of the process of ion field emission (IFE). This process becomes important for hard spectra of the incident radiation. We combine the processes of dust destruction (UV and X-ray sublimation and IFE) with the photoionization code that includes H, H$_2$ and the 12 most abundant astrophysical elements. Our code follows, both in space and in time, the fractions of metals that are depleted into gas and those that are in gaseous phase.  As dust is being destroyed, metals are gradually being recycled into gas.  Our code allows for any initial distribution of the grain sizes, and, for each grain size and wavelength, opacities are evaluated by interpolation on the grids of opacities computed by Draine \\& Lee (1984; DL in the following) and Laor \\& Draine (1993).  While on one side showing how GRBs modify their surroundings (and hence the inferred properties of the environment as measured at later times), we also show how monitoring the time-variability of several observables during the GRB event can yield powerful information on the characteristics of the environment.  In particular, in this paper, we discuss the simultaneous behaviour of the X-ray opacity and the optical opacity, and show how they depend on the size and density of the region. In two companion papers (Perna, Lazzati \\& Fiore 2002, paper II, and Lazzati \\& Perna 2002, paper III) we will be discussing, respectively, the time-dependence of the optical opacity on the distribution of dust grain sizes in the medium and the effect of time-dependent ionization on the appearance of X-ray spectra, focusing on the measure of the continuum absorption (equivalent column density, see Lazzati \\& Perna 2002) and on the appearance of enhanced metallicities (or deeper metal edges) similarly to the case of warm absorbers (Done et al. 1992; Zdziarski et al. 1995).  This paper is organized as follows: in \\S 2 we discuss the various dust destruction mechanisms and we derive (\\S 2.2) a mathematical expression to describe the evolution of the grain size due to the process of IFE.  In \\S 3 we describe how the code handles, during the process of dust sublimation, the transfer of metals from dust into gas, and how the initial dust distribution evolves in time. The initial conditions for dust content and metal depletion pattern are described in \\S 4, while in \\S 5 we discuss the photoionization and photodissociation of H$_2$ and how the code handles radiative transfer in optically thick regions. The evolution of metals is described in \\S 6, while the basics of the radiative transfer and the dust opacities are discussed in \\S 7. Our results for the evolution of the combined X-ray and optical opacities and their relevance for GRB environments are presented in \\S 8, while the observational perspectives (with both current and future instrumentation) for detection of the time-dependent effects that we describe, are discussed in \\S 9. Finally, our work is summarized in the last section, \\S 10. ",
        "conclusions": "The discovery of the bright GRB longer-wavelength counterparts, in particular in X-ray and UV, has brought up the problem of simultaneous photoionization and dust destruction under an intense radiation field. We have tackled the problem in this paper, and the principal results of our work are the following:  \\begin{itemize}  \\item We have generalized the mathematical formulation of the dust destruction mechanisms under an intense radiation field to include the process of ion field emission, and we have discussed its relative importance with respect to X-ray and UV sublimation as a function of grain size, intensity and hardness of the incident spectrum.  \\item We have developed a time-dependent radiative transfer code that simultaneously follows and combines the process of dust destruction through UV/ X-ray sublimation and through IFE with the evolution of the ionization states of metals due to photoionization and their relative radiative transitions. The code keeps track of the relative fractions of metals in gaseous and in dust-depleted phase as a function of space and time, and, as such, it follows the gradual recycling into gas of metals while dust is being destroyed.  The code allows for any initial distribution of grain sizes and any value for the dust-to-gas ratio and uses, through interpolation, the tabulated opacities computed by Draine \\& Lee (1983).  To reduce the running times of the simulations, we have developed (and provided in the Appendix) an approximate expression for the Planck averaged absorption opacity that reproduces the numerical results by Draine \\& Lee within $\\sim15\\%$ in the range of grain sizes $0.01\\le{}a_{-5}\\le10^2$ and temperatures $10\\le{}T\\le5\\times10^4$~K.  \\item We have applied our code to show how GRBs affect the inferred X-ray and optical extinction depending on the type of environment, and how monitoring these quantities {\\em during} the GRB event can provide powerful clues on the characteristics of the environment itself. In fact, we have shown how a measurement of X-ray absorption, $N_{\\rm H}$, and optical extinction, $A_V$, at a {\\em single} time can yield dust-to-gas ratios that are either lower or higher than the real value, depending on the density and size of the absorbing region, which determines the relative reduction in the high and low-energy opacities.  Prompt and continuous time monitoring of the opacities in various bands will be soon possible with {\\em Swift} and REM.  Whereas a knowledge of the type of environment in which GRBs occur is relevant for constraining GRB progenitors, the inferred metallicites and dust content are particularly important for a fair and comprehensive study of high-redshift galaxies, whose inner and denser cores are inaccessible to QSO absorption studies, and for which GRBs constitute the only probe available so far.  \\end{itemize}"
    },
    "0206/astro-ph0206359_arXiv.txt": {
        "abstract": "We consider here the class of compact, isolated, high--velocity \\hi clouds, CHVCs, which are sharply bounded in angular extent with no kinematic or spatial connection to other \\hi features down to a limiting column density of 1.5$\\times$10$^{18}$ cm$^{-2}$.  We describe the automated search algorithm developed by de\\,Heij et al.  (2002a) and applied by them to the Leiden/Dwingeloo Survey north of $\\delta= -28^\\circ$ and by Putman et al. (2002) to the Parkes HIPASS data south of $\\delta=0\\deg$, resulting in an all--sky catalog numbering 246 CHVCs.  We argue that these objects are more likely to represent a single phenomenon in a similar evolutionary state than would a sample which included any of the major HVC complexes.  Five principal observables are defined for the CHVC population: \\,(1) the spatial deployment of the objects on the sky, (2) the kinematic distribution, (3) the number distribution of observed \\hi column densities, (4) the number distribution of angular sizes, and (5) the number distribution of line widths.  We show that the spatial and kinematic deployments of the ensemble of CHVCs contain various clues regarding their characteristic distance.  These clues are not compatible with a location of the ensemble within the Galaxy proper. The deployments resemble in several regards those of the Local Group galaxies.  We describe a model testing the hypothesis that the CHVCs are a Local Group population. The agreement of the model with the data is judged by extracting the observables from simulations, in a manner consistent with the sensitivities of the observations and explicitly taking account of Galactic obscuration.  We show that models in which the CHVCs are the \\hi counterparts of dark--matter halos evolving in the Local Group potential provide a good match to the observables, if account is taken of tidal and ram--pressure disruption, the consequences of obscuration due to Galactic \\hi and of differing sensitivities and selection effects pertaining to the surveys.  A representative sample of CHVCs has been studied with high angular resolution (sub-arcminute) using the WSRT and with high \\NH~sensitivity ($<10^{17}$\\,cm$^{-2}$) using the Arecibo telescope. The picture that emerges is a nested morphology of CNM cores shielded by WNM cocoons, and plausibly surrounded by a Warm Ionized Medium halo.  These observations lead to indirect contraints on the distances, ranging from 150 to 850 kpc. ",
        "introduction": "\\hi structures found lying in the velocity regime harboring the high--velocity clouds manifest themselves in different forms; considered collectively, these clouds either do not represent a single astronomical phenomenon, or (more likely) do not represent the phenomenon at a single stage in its evolution, with all members seen under the same physical circumstances.  The variety of positional and kinematic properties may correspond to differing evolutionary histories.  The Magellanic Stream, for example, is identified as associated with the Large Magellanic Cloud because of positional and kinematic coincidences of part of the Stream with the LMC.  It is composed of material either tidally stripped from the Milky Way or from the LMC itself, or gravitationally captured from elsewhere, but in any case constrained to follow the orbit of the LMC, and thus presumably lies at the LMC distance of some 50 kpc.  There are several other major streams of HVCs, notably Complexes M, C, and A, which can also be traced over many tens of degrees and which, like the Magellanic Stream, individually show substantial spatial and kinematic coherence.  The distance of only one complex has been bracketed with any certainty: absorption--line measurements (van Woerden et al. 1999, Wakker 2001) toward Complex A place it within the range $4 < d < 10$ kpc.  The histories of the Magellanic Stream and of Complex A (and probably of the other large complexes as well) have been multifaceted, with exposure to strong radiation fields and tidal distortions by the Milky Way.  Braun \\& Burton (1999; BB99) identified CHVCs as a subset of the anomalous--velocity gas which might be characteristic of a single class of objects, whose members plausibly originated under common circumstances and share a common subsequent evolutionary history.  The discussion of the CHVCs hypothesized that the large HVC complexes (other than the Magellanic Stream) were once similar objects, but which now are relatively nearby.  We describe in \\S\\,2 the search algorithm that has been applied to the all--sky coverage afforded by recent \\hi surveys of the northern and southern skys.  In \\S\\,3 we show the principal observable quantities of the all--sky ensemble, and briefly remark on conclusions which follow directly from these observables.  In \\S\\,4, we show that a model in which the CHVCs are the \\hi counterparts of dark--matter halos evolving in the Local Group potential provides a good match to the observables, but only if the simulation is sampled as if it were being observed with the sensitivities of the available observations and if explicit account is taken of obscuration by \\hi in the Galaxy. In \\S\\,5, we show WSRT and Arecibo high--resolution imaging of a selection of individual CHVCs, stressing the importance of the synthesis observations to study of the cold cores and of the single--dish observations to study of the diffuse halos. The high--resolution imaging supports several indications of substantial distances. ",
        "conclusions": "The concept of a distinct class of compact, isolated high--velocity clouds has been objectively developed by application of a search algorithm to the LDS and HIPASS datasets, resulting in an all--sky catalog of CHVCs, high--contrast \\hi features which are at best only marginally resolved with half--degree angular resolution and which are sharply bounded, such that emission at column densities above a few times 10$^{18}$ cm$^{-2}$ is unconfused by extended emission in the source environment. We identify five principal observables from the all--sky catalog: the spatial and kinematic deployments, and the number distributions of angular sizes, \\NH, and linewidth. We show that agreement of models with the data must be judged by extracting these same obervables from the simulations in a manner consistent with the sensitivities of the observations, and explicitly taking into account the limitations imposed by obscuration due to foreground \\hi in the Milky Way.  General considerations of these observables show compatibility with characteristic distance scales of hundreds of kpc.  Since the inception of high--velocity cloud research, the possibility of an extragalactic deployment of these clouds has been critically considered as a possibility; among others, the discussions by Oort (1966, 1970, 1981), Verschuur (1969, 1975), Giovanelli (1981), Bajaja, Morras, \\& P\\\"oppel (1987), Wakker \\& van Woerden (1997), Blitz et al. (1999), and BB99 are particularly relevant here. Our simulations of a Local Group population of CHVCs have accounted explicitly for the obscuration resulting from the foreground Galactic \\hi and the observational attributes of the existing survey material. We have shown that models in which the CHVCs are the \\hi counterparts of dark--matter halos evolving in the Local Group can provide a good match to the observables, if the simulations are ``observed'' in accordance with the LDS and HIPASS parameters.  High--resolution imaging confirms the nested core/halo geometry expected if the CNM is to be stable in the presence of an ionizing radiation field of the sort expected in the Local Group environment. The cores contribute typically about 40\\% of the \\hi flux, while covering about 15\\% of the surface of the CHVC.  The imaging also allows supports two independent distance estimates.  A distance for CHVC\\,$125\\!+\\!41\\!-\\!207$ follows from assuming rough spherical symmetry and equating the well--constrained volume and column densities of the compact cores. Another distance constraint (coupled with a dark--to--visible mass ratio) follows from consideration of the stability of CHVCs having multiple cores in a common envelope but having large relative velocities.  The available evidence suggests that CHVCs have characteristic sizes of about 2 kpc, \\hi masses in the range $10^{5.5}$ to $10^7$ M$_\\odot$, and are typically seen at distances of hundreds of kpc.  If at such distances, the failure to detect stars would imply that CHVCs are very primitive proto--galactic objects dominated by dark--matter halos, plausibly the the missing Local Group satellite systems predicted by Klypin et al. (1999) and Moore et al. (1999)."
    },
    "0206/astro-ph0206390_arXiv.txt": {
        "abstract": "We examine the significance of the first metal-free stars (Pop III) for the cosmological reionization of H~I and He~II.  These stars have unusually hard spectra, with the integrated ionizing photon rates from a Pop III stellar cluster for H~I and He~II being 1.6 and $10^5$ times stronger respectively than those from a Pop II cluster.  For the currently favored cosmology, we find that Pop III stars alone can reionize H~I and He~II at redshifts, $z \\simeq$ 9 (4.7) and 5.1 (0.7) for continuous (instantaneous) modes of star formation.  More realistic scenarios involving combinations of Pop III and Pop II stellar spectra yield similar results for hydrogen. Helium never reionizes completely in these cases; the ionization fraction of He~III reaches a maximum of about 60\\% at $z \\sim 5.6$ if Pop III star formation lasts for $10^9$ yr. Future data on H~I reionization can test the amount of small-scale power available to the formation of the first objects, and provide a constraint on values of $\\sigma_8$ $\\la$ 0.7. Since current UV observations indicate an epoch of reionization for He~II at $z \\sim 3$, He~II may reionize more than once.  Measurements of the He~II Gunn-Peterson effect in the intergalactic medium at redshifts $z \\ga$ 3 may reveal the significance of Pop III stars for He~II reionization, particularly in void regions that may contain relic ionization from early Pop III stellar activity. ",
        "introduction": "The nature, formation sites, and epochs of the first stars in the universe are some of cosmology's most intriguing yet unresolved questions today. Theoretical studies of the effects of these objects on the high-redshift intergalactic medium (IGM) and on galaxy formation have a rapidly expanding literature, driven in part by the potential to test the predictions from such theories with data in the near future. Recent work on the first stars has focused on signatures such as the effects of stellar radiation and nucleosynthesis on their host galaxies and the IGM \\citep{go97,hl97, ferrara00, abia, rgshull02} -- loosely grouped under stellar and supernova (SN) feedback -- and the potential presence of large numbers of stellar remnants in galactic halos in baryonic dark matter scenarios \\citep{fields}.  These studies were directly motivated by observations of the reionization \\citep{becker, kriss} and trace metal enrichment (\\citealt{song01}, 2002) of the high-$z$ IGM, and the detection of solar-mass dark objects in our Galactic halo by microlensing experiments \\citep{alcock}. Although it is unclear if the same population of early stars can be tied unambiguously to all of these data, it is likely that they contributed significantly to the ionizing photon budget and metal production at early times.   Beyond such signatures, there has also been considerable interest in the typical masses of the first stars and the preferred environments, if any, in which they form. There is no theoretical basis on which one can {\\it a priori} rule out a stellar initial mass function (IMF) that was different in the past. Indeed, arguments for a primordial IMF biased towards higher masses have been proposed for some time \\citep{carr84,larson, abel,herferr,bromm,naka}, although there is no indication for an environment-dependent IMF from data of a variety of local star-forming regions \\citep{kroupa}. As for the nature of the typical galaxy that hosts the first stars, this depends critically on the availability of coolants within virialized halos so that the fragmentation necessary for star formation may commence. Several authors \\citep{teg97,ciardi} have argued that modest levels of early stellar activity can generate sufficient far-ultraviolet radiation in the Lyman-Werner bands (11.2 -- 13.6 eV) to photodissociate all of the remaining H$_2$ in the universe, well before the associated H~I ionizing flux has built up to values sufficient for reionization. A long pause in global star formation would then ensue, owing to ``negative feedback'', and would resume only when halos of virial temperature $\\ga$ 10$^4$ K collapse, corresponding to the threshold for the onset of H line cooling.  One way to overcome such negative feedback might be through the presence of X-rays or ionizing UV photons from the first luminous sources, which could boost the free electron fraction and hence the amount of H$^{-}$-catalyzed H$_2$, leading to a compensatory positive feedback \\citep{rgshull01}. It remains unclear how effective this is in overcoming the negative feedback from infrared and Lyman-Werner band photons \\citep{har,venkgs01}. There is also the possibility that sufficient metals are injected into the interstellar medium (ISM) soon after the very first stars form. In this case, the distinction between halos cooling by H$_2$ versus H becomes irrelevant, assuming that the metals can be retained in the cold star-forming gaseous component within individual halos. This problem remains unresolved currently, but it is clear that the chemistry of high-$z$ halos is critical to when and where the first stars formed.  {\\it A priori}, we would expect early generations of stars forming from primordial gas to be metal-free in composition, although no surviving members of such populations have been detected to date.  Recent studies of stars of zero metallicity, $Z$ (\\citealt{tumshull,bromm, cojazzi}; \\citealt{schaerer}, and references therein), which we henceforth refer to as Pop III, have demonstrated that $Z = 0$ stars are fundamentally different in nature and evolutionary properties from their low-$Z$ counterparts. In particular, \\citet{tumshull} showed, in a calculation of the zero-age main sequence of these stars, that their harder ionizing spectra could be relevant for both the H~I and He~II reionization of the IGM. This work has subsequently been extended to full calculations of the evolving spectra of Pop III stellar populations in \\citet{tsv02} (henceforth Paper I), where we found that, for a Salpeter IMF, the integrated ionizing photon rate from a Pop III cluster for H~I and He~II is respectively 1.6 and $10^5$ times stronger than that for a $Z = 0.001$ cluster \\citep{starburst} of the same mass\\footnote{In this comparison, we do not include the contribution of Wolf-Rayet stars which can boost the ionizing radiation from a Pop~II stellar cluster. In fact, the existence of the Wolf-Rayet phase in $Z=0$ stars is questionable, as discussed in Paper I, given that these objects are unlikely to experience strong mass loss. We are primarily interested here in a direct comparison of the ionizing radiation from the main-sequence phases of Pops III and II.}.  In Paper I, we examined the evolving spectra of Pop III stars and their observational signatures such as broad-band colors and emission lines. In this paper, which is intended as a companion work to Paper I, we focus on the significance of such metal-free stellar populations for cosmological reionization, under the assumption that they form in a present-day IMF.   We present a brief review here of the status of theoretical models and data on reionization.  Spectroscopic studies of high-$z$ quasars and galaxies blueward of their rest-frame H~I and He~II Ly$\\alpha$ emission have revealed that He~II reionization occurs at $z \\sim$ 3 \\citep{kriss} and that of H~I before $z \\sim 6$ \\citep{becker}. Such spectroscopic observations, along with increasingly precise data on the cosmic microwave background (CMB), are beginning to place strong complementary bounds on the redshift of H~I reionization, $z_{\\rm reion, H}$. At one end, current CMB data on the temperature anisotropy at degree and sub-degree scales provide an upper limit of about 0.3 for the electron-scattering optical depth to reionization, which may be translated into a model-dependent constraint of $z_{\\rm reion, H}$ $\\la$ 25 \\citep{wangtz01}.  Ongoing and future CMB observations (see http://www.hep.upenn.edu/$\\sim$max/cmb/\\newline experiments.html, and http://background.uchicago.edu/$\\sim$whu/cmbex.html for links to various CMB experiments) will provide improved constraints on $z_{\\rm reion, H}$ through the detection of polarization in the CMB at large angular scales \\citep{staggs}. At the other end of the range for $z_{\\rm reion, H}$, the IGM appears to be highly ionized up to $z \\sim 6$ \\citep{fan1,dey}. \\citet{becker} recently detected the H~I Gunn-Peterson (GP) trough in the spectrum of the highest-redshift quasar known to date at $z = 6.28$ \\citep{fan3}, which may indicate that H~I reionization occurs not far beyond $z \\sim 6$.  This claim has been challenged, however, by the subsequent observation of Ly$\\alpha$ emission in a $z = 6.56$ galaxy \\citep{hu02}.  The extent to which this Ly$\\alpha$ emission line has been eroded by the damping wing of Ly$\\alpha$ absorption in the IGM \\citep{escude} is, however, unclear.  It is therefore difficult to assess whether H~I reionization is complete at $z \\sim$ 6.5 based on this one object.  Although the detection of the GP trough in a single line of sight is not definitive evidence of the global reionization of the IGM, it may probe the end of the gradual process of inhomogeneous reionization, coinciding with the disappearance of the last neutral regions in the high-$z$ IGM.  Reionization of H~I at $z$ $\\sim$ 6 -- 9 would still be consistent with the lower end of the range of redshifts, $z \\sim$ 6 -- 20, predicted by theoretical models, both semi-analytic \\citep{tsb94,gs96,hl97,valsilk,mhrees99,mirees00} and based on numerical simulations \\citep{cenost93,gnedin00,ciardi,benson2}.   The reionization of the IGM subsequent to recombination at $z \\sim$ 1000 is thought to have been caused by increasing numbers of the first luminous sources. Although there are a variety of models for the astrophysical objects or processes that could have reionized the IGM, the leading scenarios involve photoionization by sources with soft or hard ionizing spectra, or equivalently stellar-type or QSO-type models respectively. Clearly, this division of source populations according to their spectral properties is no longer valid if the first stars generated hard ionizing radiation by virtue of their metal-free composition (Paper I) or if they formed in an IMF biased towards extremely high masses \\citep{bromm}.  The large majority of currently favored reionization models involve stars rather than QSOs, for a number of observationally motivated reasons (see \\citealt{venk02} for a detailed discussion on this point). These reasons include the apparent decrease in the space density of large, optically bright quasars up to $z \\sim$ 6.3 beyond a peak at $z \\sim$ 3 \\citep{fan3,shaver}, so that their UV and X-ray photons are insufficient for H~I reionization \\citep{mad99,venkgs01}. However, the nature of the reionizing sources is highly uncertain at present, and one can neither confirm nor exclude stars, faint QSOs (``mini-QSOs''), or a combination of high-$z$ source populations. In this paper, we are not specifically advocating that Pop III stars are solely responsible for reionization. Rather, our main goal is to examine the consequences of an epoch of metal-free star formation for the H~I and He~II reionization of the IGM, given a model of reionization with currently favored values for the input parameters.   We organize the paper as follows. In \\S 2, we present the newly calculated evolving spectra of Pop III from Paper I, and we describe the reionization model used in this work. In \\S 3, we present our results on the effects of Pop III and Pop II stars on the H~I and He~II reionization of the IGM for a number of potential scenarios of high-$z$ star formation. We also discuss the potential constraint offered by the reionization epoch on the amount of small-scale power in structure formation models, and we speculate on the fate of partially or fully ionized He~III in underdense regions of the IGM whose detection lies on the threshold of current capabilities.  We present our conclusions in \\S 4. ",
        "conclusions": "We showed in Paper I that Pop III stars have unusually hard spectra and elevated H~I and He~II ionizing photon rates. These properties motivated this work, where we examined the role played by these objects in H~I and He~II reionization through a semi-analytic reionization model described by a reasonable set of parameters in the currently favored cosmology. Our general conclusions are:   \\begin{itemize}  \\item[1.] We find that Pop III stars can be cosmologically significant for reionization, particularly for He~II. For Pop III alone, H~I and He~II reionize at redshifts $z_{\\rm reion, H}$ $\\sim$ 9.0 (4.7) and $z_{\\rm reion, He}$ $\\sim$ 5.1 (0.7) for continuous (bursty) modes of star formation.  \\item[2.] We also considered a more realistic scenario involving a Pop III phase of (continuous) star formation which switches to Pop II after a self-enrichment timescale for primordial star-forming gas. We find that H~I reionization occurs at $z_{\\rm reion, H}$ $\\sim$ 8.7 or 9.0, depending on whether the Pop III stage lasts $10^8$ or $10^9$ yr respectively. He~II never reionizes completely in either case, although the ionization fraction of He~III reaches a maximum of about 60\\% at $z \\sim 5.6$ for a $10^9$ yr self-enrichment timescale.   \\item[3.] Since the reionization epoch is sensitive to the power available on small scales, data on H~I reionization can critically test, and possibly rule out, low values of $\\sigma_8$ ($\\la$ 0.7), particularly if $n <$ 1.  \\end{itemize}  By measuring $z_{\\rm reion, H}$ from the CMB and high-$z$ spectroscopic studies, and by using direct imaging techniques to detect Pop III stellar clusters (Paper I), one can strongly constrain the role played by Pop III stars in H~I reionization. The current evidence for a complete H~I GP trough, and hence $z_{\\rm reion, H}$, comes from the spectrum of a single $z = 6.28$ QSO.  The acquisition of more GP data along more lines of sight to sources at $z \\sim$ 6 -- 9 is required to adequately represent how the appearance and duration of the GP trough varies with redshift.  Such observations are within the capabilities of the {\\it Sloan Digital Sky Survey}, which should detect about 20 bright quasars at $z \\ga$ 6 during the course of the survey \\citep{becker}, and are important targets in the planning of the {\\it Next Generation Space Telescope}.   The significance of Pop III stars for He~II reionization can be tested by future measurements of the He~II GP effect in the IGM at $z \\sim$ 3 -- 5, particularly in underdense regions of the IGM, which may not have had the time to recombine by $z \\sim 3$ after experiencing ionization by Pop III stars at early times. These relic ionized voids may retain the unique spectral imprint of the first, metal-free stellar populations."
    },
    "0206/astro-ph0206029_arXiv.txt": {
        "abstract": "{We revisit the nature of the far infrared (FIR)/Radio correlation by means of the most recent models for star forming galaxies, focusing in particular on the case of obscured starbursts. We model the IR emission with our population synthesis code, GRASIL (Silva et al. 1998). As for the radio emission, we revisit the simple model of Condon \\& Yin (1990). We find that a tight FIR/Radio correlation is natural when the synchrotron mechanism dominates over the inverse Compton, and the electrons cooling time is shorter than the fading time of the supernova (SN) rate. Observations indicate that both these conditions are met in star forming galaxies, from normal spirals to obscured starbursts. However, since the radio non thermal (NT) emission is delayed, deviations are expected both in the early phases of a starburst, when the radio thermal component dominates, and in the post-starburst phase, when the bulk of the NT component originates from less massive stars. We show that this delay allows the analysis of obscured starbursts with a time resolution of a few tens of Myrs, unreachable with other star formation (SF) indicators. We suggest to complement the analysis of the deviations from the FIR/Radio correlation with the radio slope (q--Radio slope diagram) to obtain characteristic parameters of the burst, e.g.\\ its intensity, age and fading time scale. The analysis of a sample of compact ULIRGs shows that they are intense but transient starbursts, to which one should not apply usual SF indicators devised for constant SF rates. We also discuss the possibility of using the q--radio slope diagram to asses the presence of obscured AGN. A firm prediction of the models is an apparent radio excess during the post-starburst phase, which seems to be typical of a class of star forming galaxies in rich cluster cores. Finally we discuss how deviations from the correlation, due to the evolutionary status of the starburst, affect the technique of photometric redshift determination widely used for the high-z sources. ",
        "introduction": "In recent years the study of starburst galaxies has become a very popular subject because of its intimate connection with the global star formation history of the Universe. On one side high redshift observations in the optical bands probe rest frame spectral regions that are highly affected by even tiny amounts of ongoing star formation and dust extinction. On the other, theoretical models following the paradigm of the hierarchical clustering scenario predict that merging induced star formation should be highly enhanced in the past. Current estimates of the star formation rate (SFR) of the Universe have thus been interpreted on the basis of our understanding of local analogous galaxies, in particular through UV continuum and optical line emission. However in local  starbursts a significant fraction of the ongoing star formation may be hidden to UV and optical estimators. In fact, though starburst galaxies were initially selected for the prominence of their optical emission lines, it appears that this criterion excludes other actively star forming objects and possibly limits our understanding to a small phase of their evolution. After the IRAS satellite, it became clear that SF is also highly enhanced in very and ultra luminous infrared galaxies that are otherwise highly attenuated in the optical. With space densities similar to those of quasars (Soifer et al. 1986) and total infrared luminosities spanning the range  10$^{11} $--10$^{12}$L$_{\\odot}$ and above 10$^{12}$L$_{\\odot}$, respectively, Luminous  and Ultraluminous Infrared galaxies (LIRGs and ULIRGs) are the most luminous objects in the local Universe. Evidence of the important role played by dust reprocessing was provided by the detection of a diffuse FIR background whose high intensity (equal to or higher than that of the optical, e.g. Hauser et al. 1998) implies that these galaxies are undergoing intense star formation activity (Puget et al. 1996; Dwek  et al. 1998). Furthermore, the advent of the Infrared Space Observatory (ISO), in combination with the availability of new ground facilities such as SCUBA on the JCMT, have discovered numerous high-z galaxies with enhanced IR emission (e.g. Elbaz et al. 1999; Smail et al. 2000; Barger et al. 2000).  Silva et al.\\ (1998) have first introduced the concept of age-selective obscuration, to explain the features of the observed spectral energy distribution (SED) of star forming galaxies, from normal spirals to dust obscured starbursts, from the UV to the sub-mm. In this model young stars are supposed to originate within molecular clouds and correspondingly their light is attenuated more than that of older stars, that already got rid of their parental cloud. The UV light in many starbursts is thus dominated by the older stars rather than by the younger populations. With the same assumptions, Granato et al.\\ (2000) reproduced the observable properties of local galaxies (in particular the IRAS luminosity function), working within the context of structure formation through hierarchical clustering, which has successfully confronted a wide range of observations on large scale structure and microwave background anisotropies. They showed that the concept of age-selective obscuration could explain the difference between the galactic {\\it extinction} law and the {\\it attenuation} law observed in starburst galaxies (Calzetti, Kinney, \\& Storchi-Bergmann 1994).  Poggianti, Bressan, \\& Franceschini (2001) have recently investigated the optical spectra of very luminous infrared galaxies to constrain the recent history of SF and the dust extinction characteristic of various stellar populations. They have found that the most plausible explanation for their unusual combination of strong H$\\delta$ absorption and moderate [OII] emission is again age-selective extinction. Indeed HII regions (wherein the [OII] emission originates) are highly embedded and thus are affected by a greater extinction compared to the older stellar populations which are responsible for the Balmer absorption. Under standard assumptions for the IMF, the SFR derived from the fit of the optical spectrum (continuum, absorption and emission lines) may account for a small fraction of the FIR emission. Moreover, even complemented with the information on the FIR flux, the optical-UV spectrum is not enough to identify univocal solutions. Further evidence along this direction is provided by recent observations of UV properties of ULIRGs (Goldader et al.\\ 2002).  These studies underline an intrinsic difficulty of evaluating the properties of massive starbursts only from their UV, optical and even NIR properties and the natural way out from this impasse seems provided by studies at longer wavelengths. The capability of FIR and radio spectral regions to reveal otherwise hidden complex phenomena in star forming galaxies, is testified by the existence of a \"miraculous\" correlation between their properties in these spectral windows. The FIR/Radio correlation is locally well established over a significant range of luminosity, from normal spirals to the most extreme ULIRGs and its small scatter states the universal proportions with which energy is radiated away at IR and radio wavelengths. In spite of its obscure nature, its utility appears in several aspects, beside being a firm tool within the manifold of star formation indicators. For instance the validity of the FIR/Radio correlation has been recently confirmed up to redshift $\\simeq 1.3$ (Garret 2001) and it is widely extrapolated much beyond, to estimate the redshift of more distant objects (e.g. Carilli \\& Yun 2000). Also, deviations from the correlation observed toward the central regions of rich cluster of galaxies, where a significant fraction of star forming galaxies show a radio excess, are used to trace the effect of the hot intracluster medium on their galactic magnetic field (Gavazzi \\& Jaffe 1986; Miller \\& Owen 2001).  So far there have been many attempts to explain the FIR/Radio correlation but all have soon or late invoked a fine tuning of the relevant physical properties, such as the intensity of the radiative and magnetic energy density (e.g.\\ Lisenfeld, Volk \\& Xu 1996). In this paper we revisit this correlation by combining our spectrophotometric code GRASIL, particularly suited to the study of the IR properties of dusty galaxies, with a new model of radio emission. The latter essentially follows the recipes by Condon (1992), but after a careful assessment of the validity of one of its basic assumptions, namely the proportionality between the non thermal (NT) radio emission and the core-collapse supernova (CCSN) rate. This fact renders the FIR/Radio correlation so robust and we provide, for the first time, a simple explanation of its universality. For the same reason we show that deviations are to be expected during and soon after the starburst episode, and we suggest that they can help in constraining the star formation history of these galaxies, something that cannot be done with optical, NIR and even FIR observations.  In Sect.~\\ref{sfir} we briefly describe our population synthesis code for dusty galaxies. Sect.~\\ref{sradio} is devoted to the new model of radio emission. In Sect.~\\ref{scalib} we describe our calibration of the NT radio emission model and obtain new relations for the SFR against radio emission  for the case of quiescent galaxies. In Sect.~\\ref{sq} we analyse infrared and radio properties of starburst galaxies. We show that the different fading times of FIR and radio emissions may be used to reach a time resolution of a few tens of Myr, which is impossible resting only on the UV-FIR. In Sect.~\\ref{squ} we introduce a new diagnostic tool, the FIR/Radio (q) vs radio spectral slope diagram, which potentially allows the determination of the evolutionary status of a starburst in absence of a good radio spectral coverage. We examine the location of an observed sample of compact ULIRGs and discuss whether this diagram may also provide a quantitative estimate of the threshold between AGN and star formation powered ULIRGs. In Sect.~\\ref{spsbt} we analyse the evolution in the post-starburst phase and suggest that radio excess is actually an indication of the occurrence of this phase rather than an environmental effect. In Sect.~\\ref{shz} we discuss the impact of these new findings on the determination of the redshift of SCUBA sources, a method that relies on the FIR/Radio correlation. Sect.~\\ref{sconc} is devoted to our conclusions. ",
        "conclusions": "\\label{sconc}  In this paper we have revisited the nature of the FIR/Radio correlation observed in star forming galaxies. To understand its origin and range of validity we have utilized models of normal star forming and starburst galaxies. The  infrared emission has been estimated with our population synthesis code, GRASIL, which is particularly suited for the prediction of the SED of star forming galaxies, from the UV to the sub-mm. As for the radio emission we have adopted a model which extends the capabilities of GRASIL into the radio regime, essentially following the prescriptions given by Condon \\& Yin (1990). Before adopting it we have reviewed all the possible sources of radio emission related to the star formation activity, with particular emphasis on the integrated properties of stellar populations.  In agreement with previous studies, we have found that the fraction of the NT emission due to radio supernova remnants is about 6\\% of the total and that other discrete sources provide a negligible contribution. Almost all the NT emission thus come from diffuse electrons possibly injected into the interstellar medium by CCSN events, and adopting a relation between the NT emission and the  CCSN rate seems the safest way to proceed. However, contrary to previous studies, we have also shown that the latter relation, which is at the base of the tightness of the FIR/Radio correlation, is a natural outcome {\\it whenever synchrotron losses happen on timescales that are shorter than the fading time of the CCSN rate}.  This is certainly the case in normal spirals, where the global SFR changes very little over the last billion years. But the situation may be very different in luminous starbursts, where the SN rate may change significantly in a timescale typical of the lifetime of the most massive stars. Indeed, previous studies claimed that the existence of a FIR/Radio correlation under such conditions requires a fine tuning between the magnetic and radiative energy density, which is difficult to explain. We have shown that this {\\it fine tuning} is not necessary. In fact the observed prevalence of synchrotron energy losses on inverse Compton against the photons of the high stellar radiation field, indicates large magnetic fields and, consequently, guarantees {\\it very short synchrotron electron lifetimes} (Condon 1992), certainly shorter than the typical lifetime of the most massive stars. The proportionality between the NT emission and the CCSN rate thus holds even in the extreme conditions found in the luminous obscured compact starbursts. This is why the FIR/Radio correlation appears so robust.  We have calibrated the NT emission/CCSN rate relation on the observed properties of our Galaxy. With these assumptions {\\it we reproduce well the FIR/Radio correlation of normal star forming galaxies}, namely q$_{1.4GHz}$=2.3. We are thus quite confident that our model is able to reproduce both {\\it the FIR and Radio emission} of star forming galaxies, with a minimum set of well calibrated parameters.  As for the starburst galaxies, with a rapidly changing SFR, we have shown that the different fading times of the FIR and Radio emission may be used to analyse in great detail the recent star formation history in these galaxies. Indeed, in the cases of M82 and ARP220, presented in Sect.~\\ref{sq}, the combination of FIR and Radio observations allows to reach a time resolution of a few tens of Myr, which is not possible based on the UV-FIR continuum properties alone, and even on the optical/NIR emission lines diagnostics, for obscured galaxies (Poggianti et al. 2001, Bressan et al.\\ 2001). We have thus analysed in greater detail the evolution of obscured starburst galaxies under different assumptions concerning the burst characteristics, challenged by the possibility of using the deviations from the FIR/Radio correlation as a diagnostic tool to infer the recent star formation history. We have compared our results with observations of a sample of compact ULIRGs, having in mind that, in these objects, it is also not clear what fraction of IR and radio emission is possibly contributed by the AGN. Compact ULIRGs show a noticeable deviation from the average FIR/Radio relation, with their  q$_{1.49{\\rm GHz}}$ being generally lager than 2.3, a few of them with values as high as 3. Taken at face values, these deviations suggest that radio emission is depressed by a large factor, relatively to normal spirals. Understanding the interplay between FIR and radio emission in these objects is thus fundamental to make reliable predictions for high redshift dust enshrouded galaxies.  Starburst models with peak SFR reaching several hundred M$_\\odot$/yr and thereafter exponentially declining, may account for the IR and radio emission of the observed ULIRGs and are able to reproduce the observed variation of the value of q$_{1.49{\\rm GHz}}$. This view is consistent with the current idea that ULIRGs and, to a larger extent high redshift dust enshrouded galaxies, are transient phenomena that nevertheless build up a significant fraction of stars and metals (Granato et al. 2001).  The introduction of a new diagnostic diagram, the q$_{1.49{\\rm GHz}}$ vs radio slope diagram, allows us to single out the effects of starburst evolution and free-free absorption. Very young star bursts display an excess of FIR emission relative to the radio emission  because the latter is initially contributed mainly by the free-free emission process. As the starburst ages, the NT contribution increases and becomes the dominant source, while the radio slope reaches the typical values observed in synchrotron emission. Free-free absorption affects the 1.49 GHz data, introducing a trend with the higher q being accompanied by the shallower slope. The estimated optical depths for free-free absorption at 1.49 GHz are between 0.5 and 1. At 8.4 GHz, free-free absorption becomes negligible and the above trend disappears. The value of q$_{8.4{\\rm GHz}}$ is a measure of the age of the starburst. However, even in the latter diagram the slope is still affected by free-free absorption. Thus we suggest that a similar diagram between 8.4 GHz and a higher frequency range would be critical for the understanding of the evolutionary status of  compact ULIRGs because, in that case, the slope, unaffected by free-free absorption, would provide an independent estimate of the age.  If ULIRGs are transient phenomena as suggested by other independent studies, then determining their SFR from conventional estimators may be a problem. They are far from being in a stationary status, the term \"average star formation\" is meaningless, and applying standard calibrations may result in a significant error and/or apparent discrepancy between the observable themselves. One should be able to reconstruct the recent history of star formation and, for this purpose, we suggest the use of the above diagram to determine the characteristic parameters of the burst first, and then the age-averaged SFR from either the FIR or radio luminosity.  Another relevant question addressed is how reliable is the use the FIR/Radio correlation to evaluate the contribution of non thermal radiation from the central active nucleus. Among the plotted data, the symbol \"M\" indicates the position of the Seyfert 1 galaxy  Mrk~231 (UGC 08058) (Thean et al. 2000). The fact that Mrk~231 is clearly distinct from the other objects and occupies a position below any starburst model, becomes particularly evident in the q$_{8.4{\\rm GHz}}$ vs slope diagram, where the effects of free-free absorption on the q ratio are minimized.  We have also shown that during the {\\it post starburst} phase, the models reach values of q significantly lower than those of quiescent spirals, with still significant FIR luminosities. This is consistent with the detection in nearby Abell clusters of a statistically significant excess of star forming galaxies with enhanced radio emission relative to the FIR (Miller \\& Owen 2001). We suggest that these low values of q are due to an evolutionary effect rather than a direct enhancement of radio emission by interaction with the intracluster medium.  Finally we have investigated on the redshift dependence of the the FIR/Radio correlation and on its validity (through the sub-mm radio index, $s^{353}_{1.4}$) to provide a photometric redshift estimate of obscured distant galaxies. The large dispersion  of $s^{353}_{1.4}$ observed among local galaxies is compatible with the evolutionary effects discussed in Sect.~\\ref{sq}. The unknown evolutionary status of the starburst renders the $s^{353}_{1.4}$ index very unreliable at almost any redshift. We thus suggest to complement the index $s^{353}_{1.4}$ with a radio slope determination, because of its tight relation with the evolutionary phase of the starburst and the its very shallow dependence on the redshift. Other uncertainties like the sub-mm slope and the presence of a significant contribution at radio wavelengths from a central AGN, obviously worsen the above picture."
    },
    "0206/astro-ph0206503_arXiv.txt": {
        "abstract": "We describe a highly unusual microlensing event, \\sc40. Unlike most standard microlensing events, this event exhibits multiple peaks in its light curve. The Einstein radius crossing time for this event is approximately one year, which is unusually long. We show that the additional peaks in the light curve can be caused by the very small value for the relative transverse velocity of the lens projected into the observer plane ($\\tilde{v} \\approx 12.5\\pm 1.1 \\kms$). Since this value is significantly less than the speed of the Earth's orbit around the Sun ($v_{\\oplus} \\sim 30 \\kms$), the motion of the Earth induces these multiple peaks in the light curve. This value for $\\tilde{v}$ is the lowest velocity so far published and we believe that this is the first multiple-peak parallax event ever observed. We also found that the event can be somewhat better fitted by a rotating binary-source model, although this is to be expected since every parallax microlensing event can be exactly reproduced by a suitable binary-source model. A face-on rotating binary-lens model was also identified, but this provides a significantly worse fit. We conclude that the most-likely cause for this multi-peak behaviour is due to parallax microlensing rather than microlensing by a binary source. However, this event may be exhibiting slight binary-source signatures in addition to these parallax-induced multiple peaks. With spectroscopic observations it is possible to test this `parallax plus binary-source' hypothesis and (in the instance that the hypothesis turns out to be correct) to simultaneously fit both models and obtain a measurement of the lens mass. Furthermore, spectroscopic observations could also supply information regarding the lens properties, possibly providing another avenue for determining the lens mass. We also investigated the nature of the blending for this event, and found that the majority of the $I$-band blending is contributed by a source roughly aligned with the lensed source. This implies that most of the $I$-band blending is caused by light from the lens or a binary companion to the source. However, in the $V$-band, there appears to be a second blended source 0.34 arcseconds away from the lensed source. Hubble Space Telescope observations will be very useful for understanding the nature of the blends. We also suggest that a radial velocity survey of all parallax events will be very useful for further constraining the lensing kinematics and understanding the origins of these events and the excess of long events toward the bulge. ",
        "introduction": "At the time of writing, more than one thousand microlensing events are known. In addition to the original goal to search for the dark matter (Paczy\\'nski 1986), these events have also developed diverse applications (see Paczy\\'nski 1996 for a review). Most of these events are well described by the standard shape (e.g., Paczy\\'nski 1986). Unfortunately, from these standard microlensing light curves, the lens distance and mass cannot be uniquely determined (see \\S2). This degeneracy is one of the major obstacles in studies of microlensing.  Fortunately, some microlensing events show deviations from the standard shape. The parallax microlensing events are one class of these exotic microlensing events (Gould 1992). These events allow one to derive the projected Einstein radius (or equivalently, the transverse velocity) in the ecliptic plane. This additional constraint partially lifts the lens degeneracy, but is not enough to uniquely determine the mass of the lensing object; to do this some other piece of information is required, such as the angular Einstein radius. This is a rare occurrence, but a striking example of this can be seen in An et al. (2002), where parallax signatures were observed in a caustic-crossing binary-lens event; this completely broke the lens-mass degeneracy and proved to be one of the first ever measurements of the microlens mass. Alcock et al. (2001a) also claim to have made a determination of the microlens mass for a different event by utilising measurements of both the parallax effect and the microlens proper motion. However, this mass determination relies on their photometric measurement of the parallax effect, which in this instance is very small and requires confirmation (this can be done by obtaining a measurement of the astrometric parallax using HST, for example). Other approaches to this degeneracy problem include resolving the components of the lensed object, which should become possible in the near-future with the availability of suitable optical long baseline interferometry (see Delplancke, G\\'orski \\& Richichi 2001). However, in the absence of further information the parallax effect can be combined with a model for the lens kinematics, which allows important constraints on the lens to be drawn. So far about 10 microlensing parallax events have been found (Alcock et al. 1995; Mao 1999; Smith, Wo\\'zniak \\& Mao 2002; Bennett et al. 2001; Mao et al. 2002; Bond et al. 2001). Three of these events are particularly interesting (Bennett et al. 2001; Mao et al. 2002) because the long Einstein radius crossing time and the kinematics imply that the lenses are very likely intervening black holes (see also Agol et al. 2002). This is particularly exciting because these black holes may be outside the gas layer of the disk, and hence have no accretion signatures for detection in any other wavelengths such as X-ray and radio. Microlensing may be the only method to provide a complete census of the massive black holes in the Milky Way.  Nearly all of the published microlensing parallax events are identified by a slight asymmetry in the light curve due to the Earth's motion around the Sun. During the systematic search of Smith et al. (2002), which analysed the 520 microlensing events published in Wo\\'zniak et al. (2001), one event has been uncovered that can be very well fitted by the parallax model. This event shows a striking multi-peak signature. Such multi-peak events have been predicted by Gould, Miralda-Escude, \\& Bahcall (1994). The purpose of this paper is to present a detailed analysis of this event. The outline of the paper is as follows: in Section 2 we present the observational data, the data reduction procedure and our method to select the parallax events from the microlensing database. In Section 3 we present the best parallax model, while in Section 4 we explore whether \\sc40 can be fitted by alternate rotating binary-source and binary-lens models. Section 5 discusses the various approaches that can be utilised to provide a measurement of (or strong constraints on) the lens mass. Finally, in Section 6, we summarize and further discuss our results. ",
        "conclusions": "In this paper we have shown that \\sc40 is a unique microlensing event which exhibits multi-peak behavior. The event was first identified in the OGLE-II early-warning system. The four season data from 1996-2000 can be reasonably fitted by a parallax model. We have also shown that the event can be fitted by a rotating binary-source model with a somewhat better $\\chi^2$. Attempts were also made to fit the data with a face-on rotating binary-lens model; the derived $\\chi^2$ is substantially worse, although we can not exclude the possibility that a better binary-lens model can be found from a more exhaustive search.  It is unlikely that the predicted difference in flux between the parallax model and the binary-source model will be discernible (see Fig. \\ref{fig:pred}). Even though photometric observations may be unable to discriminate between these two models, such observations will be able to test whether these models are feasible or not, since they both predict a significant drop in flux (approximately 0.03 magnitudes) between the end of the 2001 season and the beginning of the 2002 season. Unfortunately, the OGLE data only covers the 1996-2000 seasons, but further data has been obtained by the PLANET collaboration.  However, even though the rotating binary-source model provided a better fit than the parallax model, it was shown in \\S\\ref{bin:circ} that the improvement in $\\chi^2$ between the two models may not mean that the binary-source model is a more-likely solution than the parallax model. In fact, the parameters for this best-fit binary-source model indicate that it is probably reproducing a mirror-image of the parallax model. This is a generic degeneracy which exists when the parallax and rotating binary-source models are fit independently. We believe that the most-likely explanation for the improvement in $\\chisq$ between the parallax and binary-source models is that the event is exhibiting both parallax and binary-source behaviour simultaneously, i.e., the major variability is caused by parallax, with a slight additional variation on top of this which is caused by some binary rotation in the source (see \\S\\ref{para+bin}). This would be very difficult to fit; however, if the orbital parameters of the binary source could be constrained by spectroscopic observations, then such a fit may be possible. Simultaneously fitting the parallax and binary-source models would enable strong constraints to be put on the lens mass, and may even provide a direct measurement of it. These spectroscopic observations may also confirm the `pure' binary-source model, in the extremely unlikely case that this is proved to be correct. It was shown in \\S\\ref{bin:circ} that the amplitude of the radial velocity variations predicted by this binary-source model should be approximately $5\\rm{km s}^{-1}$ with a period of one year, and therefore this prediction can easily be discounted by a couple of measurements with $1\\rm{km s}^{-1}$ precision. However, as was noted in \\S\\ref{fss}, these predictions from the `pure' binary-source model may violate the limit placed on $\\thetaE$ by finite source size considerations, providing further evidence that this `pure' binary-source model may not be a true physical solution.  One striking prediction of the parallax model is the projected velocity $\\tilde{v}$, which is the lowest seen in any lens candidate. Smith et al. (2002) pointed out that such low-transverse velocities may be caused by a disk-disk lensing event. However, this does not appear to be the case for \\sc40, as the source seems to be a red-clump star in the bulge (see Fig. \\ref{fig:cmd}; see also Bennett et al. 2001 for other black hole candidates). Usually, a low transverse velocity is produced when the lens and source move more or less radially (Mao \\& Paczy\\'nski 1996). Therefore it will be very important to obtain the radial velocities of \\sc40. A spectroscopic survey of all parallax events, including \\sc40, will be a worthwhile effort since radial velocities will provide further information on the source kinematics. Such a survey will also shed light on the related (unsolved) problem of the nature of the excess of long events, which was first noticed by Han \\& Gould (1996).  We have investigated the photometry and astrometry of the blended sources. The $V$-band and $I$-band data together imply that there are two blended sources within the seeing disk of the lensed source, one is aligned with the lensed source while the other is mis-aligned by about 0.83 pixels ($\\sim 0.37\\arcsec$). This mis-aligned blend could be an unrelated star or, if the lens is close enough to the observer, this could be a wide-separation companion to the lens (although if the lens is at a distance of 1.3kpc, as was suggested in \\S\\ref{blend}, then the separation would be around 400 au, which would be too great a separation for this companion to have any influence on the microlensing light curve). However, the OGLE-II data are not sufficient to decipher these two components. The Advanced Camera for Surveys on HST is the ideal instrument to resolve such components. Furthermore, to understand the nature of the aligned blending, it will be very useful to obtain spectra from the ground. If the aligned source is from a binary source companion, then spectroscopy will reveal period shifts in the radial velocity. If the light is from the lens, then cross-correlations of the spectra may still show evidence for the lens due to the relative difference between the radial velocities of the source and the lens (Mao, Reetz \\& Lennon 1998). These observations may be particularly valuable for further understanding the lensed system, including a potential determination of the lens mass.  It would be interesting to know how frequently these multiple peak events occur, if indeed they are caused by the parallax effect. Since an especially low velocity is required to produce such multiple peaks, one can gain an indication of their prevalence by studying theoretical distributions of the projected velocity, $\\tilde{v}$ (e.g., Han \\& Gould 1995). The distributions produced by Han and Gould (1995) suggest that it is extremely unlikely such events could be caused by bulge-bulge lensing, since in this instance values of $\\tilde{v} < 100 \\kms$ are strongly disfavoured. For lenses residing in the disk, on the other hand, values of $\\tilde{v}$ are expected to be much smaller, although velocities comparable to \\sc40 also appear to be highly unlikely according to their findings. One would need to perform detailed Monte Carlo simulations to gain a firmer understanding of the prevalence of such multiple peak events, and this is something which we are currently undertaking. This is important because such events can easily be overlooked when compiling microlensing catalogues since their variability differs significantly from the standard microlensing light curve and also because they take an unusually long time to reach a constant baseline (due to the exceptionally small value of $\\tilde{v}$ that is required to produce such multiple peaks). For example, the variability for \\sc40 is expected to last for well-over 5 years in total. Therefore, if the observations had started once the event was underway, this event could have easily been mistaken for a variable source.  So far at least three convincing parallax candidates have been found among the 520 microlensing events in the catalogue of Wo\\'zniak et al. (2001). These include \\sc40, along with two events that have been studied in previous papers, namely sc33\\_4505 (Smith et al. 2002) and OGLE-1999-BUL-32 (Mao et al. 2002). So it appears that the parallax rate is about 1\\%. However, as discussed in Smith et al. (2002), there are a number of marginal microlensing events that were uncovered. Detailed Monte Carlo simulations are needed to check whether a similar rate is expected in the OGLE experiments. We are currently performing such a study, the results of which will be reported elsewhere."
    },
    "0206/astro-ph0206053_arXiv.txt": {
        "abstract": "We report the first results of an ongoing survey at 22\\,GHz with the 100-m Effelsberg telescope to search for water maser emission in bright IRAS sources. We have detected water vapor emission in IC~342. The maser, associated with a star forming region $\\sim$10--15\\arcsec\\ west of the nucleus, consists of a single 0.5\\,km\\,s$^{-1}$ wide feature and reaches an isotropic luminosity of 10$^{-2}$\\,L$_{\\odot}$ ($D$=1.8\\,Mpc). Our detection raises the detection rate among northern galaxies with IRAS point source fluxes $S_{\\rm 100\\mu m}$ $>$ 50\\,Jy to 16\\%. ",
        "introduction": "To date, luminous extragalactic H$_2$O masers can be grouped into three classes: those tracing accretion disks (e.g.\\ NGC~4258); those in which the emission is either the result of an interaction between the radio jet and an encroaching molecular cloud or an accidental overlap, along the line-of-sight, between a warm dense molecular cloud and the radio continuum of the jet (e.g.\\ NGC~1052 and Mrk~348); those related to prominent sites of star formation, such as the ones observed in M~33 (the earliest known extragalactic H$_2$O masers), and later also in IC~10.  Extragalactic H$_2$O masers are preferentially detected in nearby galaxies that are bright in the infrared (\\cite{braatz97}). While nuclear masers are of obvious interest, non-nuclear masers are also important for a number of reasons: these sources allow us to pinpoint sites of massive star formation, to measure the velocity vectors of these regions through VLBI proper motion studies, and to determine true distances through complementary measurements of proper motion and radial velocity (e.g. \\cite{greenhill93}).  We have therefore observed the nearby spiral galaxies IC~342 and Maffei~2, both of which exhibit prominent nuclear bars and strong molecular, infrared, and radio continuum emission. In the following we report the results of our observations. ",
        "conclusions": ""
    },
    "0206/astro-ph0206265_arXiv.txt": {
        "abstract": "High resolution, non-oscillatory, central difference (NOCD) numerical schemes are introduced as alternatives to more traditional artificial viscosity (AV) and Godunov methods for solving the fully general relativistic hydrodynamics equations. These new approaches provide the advantages of Godunov methods in capturing ultra-relativistic flows without the cost and complication of Riemann solvers, and the advantages of AV methods in their speed, ease of implementation, and general applicability without explicitly using artificial viscosity for shock capturing. Shock tube, wall shock, and dust accretion tests, all with adiabatic equations of state, are presented and compared against equivalent solutions from both AV and Godunov based codes. In the process we address the accuracy of time-explicit NOCD and AV methods over a wide range of Lorentz factors. ",
        "introduction": "\\label{sec:intro}  The earliest attempts at simulating relativistic flows in the presence of strong gravitational fields are attributed to May and White (1966, 1967) who investigated gravitational collapse in a one dimensional Lagrangian code using artificial viscosity (AV) methods \\citep{Neumann50} to capture shock waves. Wilson (1972, 1979) subsequently introduced an alternative Eulerian coordinate approach in multi-dimensional calculations, using traditional finite difference upwind methods and artificial viscosity for shock capturing. Since these earliest studies, AV methods have continued to be developed in their popularity and applied to a variety of problems due largely to their general robustness \\citep{HSW84_1,HSW84_2,CW84,Anninos98}. These methods are also computationally cheap, easy to implement, and easily adaptable to multi-physics applications. However, it has been demonstrated that problems involving high Lorentz factors (greater than a few) are particularly sensitive to different implementations of the viscosity terms, and can result in large numerical errors if solved using time explicit methods \\citep{Norman86}.  Significant progress has been made in recent years to take advantage of the conservational form of the hydrodynamics system of equations to apply Godunov-type methods and approximate Riemann solvers to simulate ultra-relativistic flows \\citep{Eulderink95,Banyuls97,Font00}. Although Godunov-based schemes are accepted as more accurate alternatives to AV methods, especially in the limit of high Lorentz factors, they are not infallible and should generally be used with caution. They may produce unexpected results in certain cases that can be overcome only with specialized fixes or by adding additional dissipation. A few known examples include the admittance of expansion shocks, negative internal energies in kinematically dominated flows, `carbuncle' effect in high Mach number bow shocks, kinked Mach stems, and odd/even decoupling in mesh-aligned shocks \\citep{Quirk94}. Godunov methods, whether they solve the Riemann problem exactly or approximately, are also computationally much more expensive than their simpler AV counterparts, and more difficult to extend the system of equations to include additional physics.  Hence we have undertaken this current study to explore an alternative approach of using high resolution, non-oscillatory, central difference (NOCD) methods \\citep{Jiang98a,Jiang98b} to solve the relativistic hydrodynamics equations. These new schemes combine the speed, efficiency, and flexibility of AV methods with the advantages of the potentially more accurate conservative formulation approach of Godunov methods, but without the cost and complication of Riemann solvers and flux splitting.  The NOCD methods are implemented as part of a new code we developed called Cosmos, and designed for fully general relativistic problems. Cosmos is a collection of massively parallel, multi-dimensional, multi-physics solvers applicable to both Newtonian and general relativistic systems, and currently includes five different computational fluid dynamics (CFD) methods, equilibrium and non-equilibrium primordial chemistry, photoionization, radiative cooling, radiation flux-limited diffusion, radiation pressure, scalar fields, Newtonian external and self gravity, arbitrary spacetime geometries, and viscous stress. The five hydrodynamics methods include a Godunov (TVD) solver for Newtonian flows, two artificial viscosity codes for general relativistic systems (differentiated by mesh or variable centering type: staggered versus zone-centered), and two relativistic methods based on non-oscillatory central difference schemes (differentiated also by the mesh type: staggered versus centered in time and space). The emphasis in the following sections is to review our particular implementations of the AV and NOCD methods and compare results of various shock wave and accretion test calculations with other published results. We also explore the accuracy of both AV and NOCD methods in simulating ultra-relativistic shocks over a wide range of Lorentz factors. ",
        "conclusions": "\\label{sec:conclusion}  We have developed new artificial viscosity and non-oscillatory central difference numerical hydrodynamics schemes as integral components of the Cosmos code framework for performing fully general relativistic calculations of strong field flows. These methods have been discussed at length here and compared also with published state-of-the-art Godunov methods on their abilities to model shock tube, wall shock and black hole accretion problems. We find that for shock tube problems at moderate to high boost factors, with velocities up to $V \\sim 0.99$ and limited only by grid resolution, internal energy formulations using artificial viscosity methods compare quite favorably with total energy schemes such as the NOCD methods, the Godunov methods using the Marquina, Roe, or Flux-split approximate Riemann solvers, and the piecewise parabolic method with an exact Riemann solver. However, AV methods can be somewhat sensitive to parameters (e.g., viscosity coefficients, Courant factor, etc.) and generally suspect at high boost factors ($V > 0.95$) in the wall shock problems we have considered here. On the other hand, NOCD methods can easily be extended to ultra-relativistic velocities ($1-V < 10^{-11}$) for the same wall shock tests, and are comparable in accuracy over the entire range of velocities we have simulated to the more standard but complicated Riemann solver codes. NOCD schemes thus provide a robust new alternative to simulating relativistic hydrodynamical flows since they offer the same advantages of Godunov methods in capturing ultra-relativistic flows but without the cost and complication of Riemann solvers or flux splitting. They also provide all the advantages of AV methods in their speed, ease of implementation, and general applicability (including straightforward extensions to more general equations of state) without explicitly using artificial viscosity for shock capturing."
    },
    "0206/astro-ph0206115_arXiv.txt": {
        "abstract": "{ We present time series observations of intermediate mass PMS stars belonging to the young star cluster IC 348. The new data reveal that a young member of the cluster, H254, undergoes periodic light variations with $\\delta$ Scuti-like characteristics. This occurrence provides an unambiguous evidence confirming the prediction that intermediate-mass pre-main sequence (PMS) stars should experience this transient instability during their approach to the main-sequence. \\par On the basis of the measured frequency $f=$7.406~day$^{-1}$, we are able to constrain the intrinsic stellar parameters of H254 by means of linear, non adiabatic, radial pulsation models. The range of the resulting luminosity and effective temperature permitted by the models is narrower than the observational values. In particular, the pulsation analysis allows to derive an independent estimate of the distance to IC 348 of about 320 pc. Further observations could either confirm the monoperiodic nature of H254 or reveal the presence of other frequencies. ",
        "introduction": "IC~348 is a young ($\\le$ 10 Myr), nearby (d$\\approx$300 pc) cluster within the Perseus complex. The cluster belongs to the Per~OB2 association, and is located near the tip of the Perseus ridge which contains other star-forming regions, such as NGC~1333 (Blaauw 1952; Herbig 1998).  A number of T Tauri stars were discovered in IC~348 by Herbig (1954) who suggested that these stars could be as young as the OB association.  The age of IC~348 has been debated in the literature due to the discrepancy between the kinematic age (1--1.4 Myr, Herbig 1998 and references therein) and that obtained from evolutionary considerations (3-20 Myr, Strom, Strom \\& Carrasco 1974; Trullols \\& Jordi 1997). According to Herbig (1998), the mean age of IC 348 inferred from its faint members is much smaller than previous evolutionary estimates, but the age spread of individual stars encompass most of evaluations in the literature.  In fact, reconstruction of the history of the cluster indicates that star formation increased dramatically about $3\\times 10^6$~years ago, with an $e$-folding time of the accelerating phase of just $1\\times 10^6$~yr (Palla \\& Stahler 2001). Current estimates of the distance to IC 348 range from 240-260 pc (Trullos \\& Jordi 1997; \\u{C}ernis 1993) to 316 pc (Herbig 1998). However, there are indications in favor of the larger distance (see discussion in Herbig 1998).  Recently, about 50 new variable stars have been discovered in IC~348 by Herbst et al. (2000) who studied their long term behavior over a period of several months. All these variable stars are classical or weak-lined T Tauri stars, while none of the early-type members showed any variability. Unfortunately, this extensive work does not provide information on the possible existence of $\\delta$ Scuti-like oscillations with time scales of the order of hours or less in the cluster members.  During the last few years a growing interest has developed in the study of the pulsational properties of young stars of intermediate mass (e.g., Marconi et al. 2001, 2002; Kurtz \\& Catala 2001 and references therein).  After the initial identification of pre--main-sequence (PMS) $\\delta$ Scuti candidates in the open cluster NGC~2264 (Breger 1972), several studies have reported results on the search for this kind of variability in known Herbig Ae stars (Kurtz \\& Marang 1995; Donati et al. 1997; Kurtz \\& Muller 1999, 2001; Marconi et al.  2000, 2001; Pigulski et al. 2000; Pinheiro et al. 2002).  These investigations were stimulated by the theoretical study of Marconi \\& Palla (1998) that established the location of the instability strip for PMS objects on the basis of nonlinear convective hydrodynamical models.  As shown in our previous investigations (Marconi et al. 2000, 2001, 2002), the comparison between observed periodicities and the predictions of linear non-adiabatic models provides useful constraints on the occurrence of radial pulsations, as well as on the intrinsic stellar parameters.  Moreover, the morphology of the PMS and post-MS tracks together with the comparison of the predicted instability strip helps to constrain the evolutionary state and the modal stability.  However, apart from Breger's candidates in NGC~2264 for which only few hours of observations are available, the other identified PMS $\\delta$ Scuti are isolated Herbig Ae stars. Obviously, the theoretical analysis outlined above is expected to give the best results for stars in clusters where the uncertainties in distance, age and reddening are less and the same for all the members. For this reason, we have selected IC~348 whose small distance, young age, and relatively large population offer a good opportunity to investigate the short time scale pulsation of some of its members.  The extensive study of Luhman et al. (1998) provides a detailed census of the stellar members of the cluster. From the comparison between the stellar parameters and the topology of the instability strip calculated by Marconi \\& Palla (1998), we have selected three PMS $\\delta$ Scuti candidates: H83, H254, and H261 (following Herbig's 1998 notation). Their properties are listed in Table~\\ref{luhman}. In this paper, we present the results of an extensive observational search that resulted in the detection of variability in one star, H254.  The paper is organized as follows: in Section 2 observations and data analysis are discussed, whereas the frequency analysis for the identified PMS $\\delta$ Scuti candidate is presented in Section 3. In Sect. 4 we compare the resulting periodicity with the predictions of PMS evolutionary and pulsation models. Some final remarks close the paper.  \\begin{table*} \\caption[]{Properties of the stars H83, H254 and H261 belonging to IC348. Values for $T_{\\rm eff}$ and $L_{\\rm bol}$ are taken from Luhman et al. (1998). For the errors in $T_{\\rm eff}$ and $L_{\\rm bol}$ see discussion in Sect. 4. \\label{luhman}} \\begin{tabular}{lcccccc} \\hline \\noalign{\\smallskip} Star  & $\\alpha$ & $\\delta$ & ST & V     & $T_{\\rm eff}$ & $L_{\\rm bol}$  \\\\ &  (J2000) &  (J2000) &    & (mag) & (K)           & $(L_{\\odot})$  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} H83  & 3 44 19.12 & +32 09 30.8 & F0 & 11.9 & $7200\\pm 170$ & $6.6\\pm 2.3$  \\\\ H254 & 3 44 31.21 & +32 06 22.1 & F0 & 10.6 & $7200\\pm 170$ & $31.4\\pm 10.1$ \\\\ H261 & 3 44 24.67 & +32 10 14.4 & F2 & 11.6 & $6890\\pm 160$ & $15.0\\pm 4.8$ \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "The present observations have revealed that the PMS F-type star H254, member of IC~348, undergoes periodic light variations with $\\delta$ Scuti-like frequency. Since IC 348 has an estimated age of less than $\\sim$10 Myr and is still actively forming stars, the discovery of small amplitude pulsations in H254 confirms the prediction by Marconi \\& Palla (1998) that intermediate-mass PMS stars should experience this transient instability during their approach to the main-sequence. Other cases presented in the literature have some ambiguity regarding the actual evolutionary state, i.e. pre--main- vs. post--main-sequence (e.g., V351 Ori described in Marconi et al. 2001). Thus, IC~348 and NGC~2264 are the only two young clusters where $\\delta$ Scuti-like  pulsations in PMS stars have been detected so far. Similar searches should be conducted on other clusters to enlarge the sample of pulsating intermediate-mass stars. Good candidates can be found in, e.g., NGC~2362 (Moitinho et al. 2001) and the Upper Sco-Cen association (Preibisch \\& Zinnecker 1999) that contain a rather large population of stars of the appropriate spectral types.  In H254, we find only one mode of pulsation with $f_4=$7.406~day$^{-1}$ (or 3.24 hr). Non-adiabatic linear models with the observed parameters of H254 show unstable modes of low order, namely the fundamental or the first overtone mode. The occurrence of just one mode of pulsations could be intrinsic or due to the detection threshold. Theoretical models of classical $\\delta $ Scuti stars predict many unstable modes, in excess of what is actually observed (e.g. Bradley \\& Guzik 2000). On the other hand, monoperiodic radially pulsating $\\delta$ Scuti stars are known (20~CVn being the best example, Chadid et al. 2001), and H254 may represent the young counterpart of this (admittedly limited) class. Spectroscopic observations of this star should be carried out to confirm the photometric period found by us and to verify whether other modes are also present. Of course, the low amplitude of the modes may render the detection quite difficult. However, the identification of a young, monoperiodic pulsating star will be very useful to shed light on the physical mechanism that limits the amplitude of the pulsations. Thus, we encourage other groups interested in stellar pulsations to consider H254 for further study."
    },
    "0206/astro-ph0206323_arXiv.txt": {
        "abstract": "In an effort to more precisely define the spatial distribution of Galactic field stars, we present an analysis of the photometric parallaxes of stars in seven Kapteyn Selected Areas. Our photometry database covers $\\sim14.9$ square degrees and contain over 130,000 stars, of which approximately 70,000 are in a color range ($0.4 \\leq R-I \\leq 1.5$) for which relatively unambiguous photometric parallaxes can be derived.  We discuss our photometry pipeline, our method of determining photometric parallaxes and our analysis efforts.  We also address the affects of Malmquist Bias, subgiant/giant contamination, metallicity and binary stars upon the derived density laws.  The affect of binary stars is the most significant of these biases -- a binary star fraction of 50\\% could result in derived scale heights that are 80\\% of the actual values.  We find that while the disk-like populations of the Milky Way are easily constrained in a simultaneous analysis of all seven fields, no good simultaneous solution for the halo is found. We have applied halo density laws taken from other studies and find that the Besan\\c{c}on flattened power law halo model ($\\frac{c}{a}=0.6, \\rho \\propto r^{-2.75}$) produces the best fit to our data. With this halo, the thick disk has a scale height of 750 pc with an 8.5\\% normalization to the old disk.  The old disk scale height is $\\sim$ 280-300 pc for our early type ($5.8 \\leq M_R < 6.8$) dwarfs and rises to $\\sim$ 350 pc for our late type ($8.8 \\leq M_R \\leq 10.2$) dwarf stars.  Corrected for a binary fraction of 50\\%, these scale heights are 940 pc and 350-375 pc, respectively.  Even with this model, there are systematic discrepancies between the observed and predicted density distributions -- discrepancies only apparent at the faint magnitudes reached by our survey. Specifically, our model produces density overpredictions in the inner Galaxy and density underpredictions in the outer Galaxy.  A possible escape from this dilemma is offered by modeling the stellar halo as a two-component system, as favored by studies of BHB/RR Lyrae stars and recent analyses of the kinematics of metal-poor stars.  In this paradigm, the halo has a flattened inner distribution and a roughly spherical, but substructured outer distribution.  Further reconciliation could be provided by a flared thick disk, a structure consistent with a merger origin for that population. ",
        "introduction": "The present stellar content of the Milky Way is a fossil record of its formation and evolution. When a star is formed, its kinematical and chemical properties reflect the state of the Galaxy at that location at the time of the star's formation.  In high density chaotic regions of the Galaxy -- such as the Galactic midplane -- that information can be quickly scrambled.  In the more remote regions of the Galaxy, the story told by ancient stars is still legible. Since lower-mass stars have lifetimes of order a Hubble time or greater, they remain in the Galaxy as echoes of the its distant past.  The most sensible and efficient strategy for reading the messages contained in old stars is to ascertain the bulk properties of stars grouped together by similar characteristics, i.e. stellar populations.  The division of Galactic field stars into distinct populations was greatly clarified by Baade's (1944) division of stars into Population I and Population II.  This system was expanded into five populations in the seminal 1957 Vatican conference (O'Connell 1958).  Even after nearly six decades of work, however, there is still some uncertainty on the exact characteristics of each population and, more importantly, what those characteristics tell us about the evolution of the Galaxy.  Indeed, there is not even a consensus on {\\it how many} populations are required to fully describe the Milky Way and whether distinct populations remain a sensible paradigm (see reviews in Majewski 1993, 1999a).  This series' first contribution to the debate (Reid \\& Majewski 1993, hereafter Paper I) used photographic starcounts data to derive an interim model for the spatial distribution of field stars toward the Galactic poles. Our second (Reid et al. 1996, hereafter Paper II) investigated small but extremely deep data sets to apply constraints to the halo luminosity function.  A later contribution (Majewski et al. 1997, hereafter Paper III) investigated an anomalous starcounts signature discovered in starcount data not presented here, and explored the possibility of contamination of our starcount data by streams of stars tidally stripped from the Sagittarius dSph galaxy.  The present discussion is a more complete and sophisticated investigation of a number of Kapteyn Selected Areas using large area, CCD-based datasets. \\S2 details the general characteristics of this dataset, while \\S3 and \\S4 detail the photometry pipeline and object classification methods, respectively.  We then use a subset of this data and the method of photometric parallax to apply stronger global constraints to the spatial distribution of Galactic field stars than any previous starcounts survey has been capable of.  \\S5 details our method of photometric parallax, \\S6-8 cover the analysis of the starcounts and our attempt to fit a self-consistent model to the data in all seven fields and \\S9 discusses the implication of those results.  \\S10 summarizes the primary results of this endeavor. ",
        "conclusions": "Our photometric parallax survey of seven Kapteyn Selected Areas has produced a number of intriguing results:  $\\bullet$ The thin disk is well described by a double exponential of raw scale height 280 pc and scale length 2-2.5 kpc.  Its axial ratio is consistent with the ratios observed in edge-on Sc galaxies.  The faintest stars in the sample show some evidence of an elevated scale height (350 pc).  Interestingly, recent results indicate that the scale height of thin disk white dwarfs may also be much higher than the canonical thin disk scale height (Majewski \\& Siegel 2002; Nelson et al. 2002).  However, if the binary fraction declines for faint stars, this would produce a similar effect.  $\\bullet$ The thick disk is well described by a double exponential of raw scale height 700-1000 pc and scale length 3-4 kpc.  Approximately 6-10\\% of the local old stars are part of the thick disk.  There is some evidence that the disk is flared, which is consistent with a kinematic heating origin for the thick disk.  $\\bullet$ If the binary fraction of Population II stars is 50\\%, the scale heights of the thin and thick disk for blue stars would increase to 350 and 900-1200 pc, respectively.  The elevated thin disk scale height of the faintest stars is 440 pc.  $\\bullet$ Fits to the halo density law do not converge in our simulations. Flattened ($\\frac{c}{a}\\sim0.6$) power law ($\\rho \\propto R_{gc}^{-2.5}$) halos are generally favored.  $\\bullet$ Some of the remaining discrepancies between our model and the data could be resolved by a dual halo model in which the inner halo is flattened and the outer halo is a roughly spherical ``can of worms\", consisting of distinct overlapping streams of stars. Such a dual halo would also reconcile many of the discrepancies in the literature into a single model.  $\\bullet$ Evidence for a flared thick disk and a dual halo would support the notion that the Milky Way has grown at least in part through the accretion of external systems.  Future contributions from this series will explore the new CCD starcount data in greater detail with the goal of strengthening our understanding of Galactic stellar populations and the underlying chemodynamical history that they reflect."
    },
    "0206/astro-ph0206378_arXiv.txt": {
        "abstract": "{ We have identified two late M dwarfs within 10\\,pc of the Sun, by cross-correlating the Luyten NLTT catalogue of stars with proper motions larger than 0.18\\,arcsec/yr, with objects lacking optical identification in the 2MASS data base. The 2MASS photometry was then combined with improved optical photometry obtained from the SuperCOSMOS Sky Surveys. The two objects (LP\\,775-\\,31 and LP\\,655-\\,48) have extremely red optical-to-infrared colours (\\colour{R}{K$_s$}$\\sim$\\magap{7}) and very bright infrared magnitudes (\\filter{K$_s$}$<$\\magap{10}): follow-up optical spectroscopy with the ESO 3.6-m telescope gave spectral types of M8.0 and M7.5 dwarfs, respectively. Comparison of their near-infrared magnitudes with the absolute magnitudes of known M8 and M7.5 dwarfs with measured trigonometric parallaxes yields spectroscopic distance estimates of $6.4\\pm1.4$\\,pc and $8.0\\pm1.6$\\,pc for LP\\,775-\\,31 and LP\\,655-\\,48, respectively. In contrast, Cruz \\& Reid (2002) recently determined spectral types of M6 for both objects, and commensurately larger distances of $11.3\\pm 1.3$\\,pc and $15.3\\pm 2.6$\\,pc. LP\\,655-\\,48 is also a bright X-ray source (1RXS~J044022.8-053020). With only a few late M~dwarfs previously known within 10\\,pc, these two objects represent an important addition to the census of the Solar neighbourhood. ",
        "introduction": "The census of the solar neighbourhood remains very incomplete. From the statistics of the Catalogue of Nearby Stars (Gliese \\& Jahrei\\ss{} \\cite{gliese91}), one can infer that about 30\\% of all stars (and probably an even larger fraction of the substellar brown dwarfs) within 10\\,pc are as yet undetected. These sources are important however, as the detailed observation of very nearby stars and brown dwarfs is one of the main starting points for investigations of the star formation process, the stellar and substellar luminosity function, and the initial mass function. Furthermore, future missions for the detection of extrasolar planets (\\eg{} SIM, TPF, Darwin) will concentrate on very nearby stars in order to reveal Earth-like planets. As a result, several major efforts are underway to provide a more detailed census of the solar neighbourhood, including the RECONS and NStars projects, out to 10 and 25\\,pc, respectively (www.chara.gsu.edu/RECONS and nstars.arc.nasa.gov). Recent discoveries in the immediate solar neighbourhood ($<$\\,15\\,pc) include early and mid M~dwarfs (Jahrei{\\ss} \\etal{} \\cite{jahreiss01}; Scholz \\etal{} \\cite{scholz01}; Phan-Bao \\etal{} \\cite{phanbao01}; Scholz \\etal{} \\cite{scholz02a}; Reid \\& Cruz \\cite{reid02}; Reid \\etal{} \\cite{reid02a}; Cruz \\& Reid \\cite{cruz02}; Reyl\\'e \\etal{} \\cite{reyle02}), late~M  and L~dwarfs (Gizis \\etal{} \\cite{gizis00}; Kirkpatrick \\etal{} \\cite{kirkpatrick00}; Reid \\etal{} \\cite{reid00}; Delfosse \\etal{} \\cite{delfosse01}; Reid \\& Cruz \\cite{reid02}; Reid \\etal{} \\cite{reid02a}; Cruz \\& Reid \\cite{cruz02}; Lodieu \\etal{} \\cite{lodieu02}), and cool white dwarfs (Reid \\etal{} \\cite{reid01}; Scholz \\etal{} \\cite{scholz02b}).  All of these objects show a proper motion larger than the NLTT (New Luyten Two Tenths catalogue; Luyten \\cite{luyten7980}) limit of 0.18\\,arcsec/yr, and indeed, many of the early and mid M~dwarfs were selected from the NLTT\\@. More than 20 years after its publication, the NLTT is still an important source in the search for nearby stars and by far not yet fully exploited. North of $\\delta$=$-33\\degree$, the NLTT catalogue is based on Palomar Schmidt plates and has a limiting magnitude of $m_{\\rm pg}$$\\sim$\\magnit{20}{5} (see, \\eg{} Reid \\& Cruz \\cite{reid02}). South of $\\delta$=$-33\\degree$, the limiting magnitude of the NLTT is only $m_{\\rm pg}$$\\sim$\\magnit{15}{5}, and thus new high proper motion surveys have been started using UK Schmidt Telescope (UKST) and ESO Schmidt plates (Scholz \\etal{} \\cite{scholz00}; Ruiz \\etal{} \\cite{ruiz01}).  Many recent additions to the census of late M~and early L~dwarfs within 15\\,pc were found in the near-infrared sky surveys 2MASS (Gizis \\etal{} \\cite{gizis00}; Kirkpatrick \\etal{} \\cite{kirkpatrick00}; Reid \\etal{} \\cite{reid00}) and DENIS (Delfosse \\etal{} \\cite{delfosse01}). Nevertheless, all these objects could in principle have been selected directly from the NLTT and/or modern Schmidt plate measurements, and in fact, all of them were subsequently identified as high proper motion objects in the NLTT or from Schmidt plates.  With the recent discovery of three nearby L~dwarfs (one with a spectroscopic distance estimate of less than 15\\,pc) in the southern sky (Lodieu \\etal{} \\cite{lodieu02}), we have demonstrated that high proper motion surveys can successfully contribute to the completion of the census of the solar neighbourhood at the stellar/substellar boundary. Furthermore, the combination of old and new proper motion catalogues with modern near-infrared sky surveys is a very effective tool in the search for nearby red dwarfs. The NLTT and 2MASS have been used to discover an M6.5 dwarf at $\\sim$\\,6\\,pc (Scholz \\etal{} \\cite{scholz01}), and five new sources were identified within 15\\,pc (and one closer than 10\\,pc) from a new proper motion survey combined with DENIS data (Reyl\\'e \\etal{} \\cite{reyle02}). Most recently, Reid \\& Cruz (\\cite{reid02}), Reid \\etal{} (\\cite{reid02a}), and Cruz \\& Reid (\\cite{cruz02}) have reported the discovery of many additions to the immediate solar neighbourhood from NLTT stars identified in the 2MASS, including five objects within 10\\,pc,  In this paper, we describe the discovery of two very nearby ($<10$\\,pc) late M~dwarfs in the southern sky, also found in a combined search of the NLTT and 2MASS\\@. The same two objects were also recently identified by Cruz \\& Reid (2002) with different results, as described below. ",
        "conclusions": "By combining data from the NLTT, 2MASS, and SSS, we have discovered two very nearby late M~dwarfs: LP\\,775-\\,31 has a spectral type of M8 and lies at a distance of $6.4\\pm1.4$\\,pc, while LP\\,655-\\,48 has a spectral type of M7.5 and lies at $8.0\\pm1.6$\\,pc. Our spectroscopic distance estimates are roughly half those recently published by Cruz \\& Reid (2002), namely $11.3\\pm 1.3$\\,pc and $15.3\\pm 2.6$\\,pc, and need to be confirmed by checking the objects for a possible binary nature and by determining trigonometric parallaxes. If confirmed, they would be among the nearest stars to the Sun and new benchmark M8 and M7.5 dwarfs, allowing detailed follow-up investigations."
    },
    "0206/astro-ph0206187_arXiv.txt": {
        "abstract": "We present results from {\\em XMM-Newton} Reflection Grating Spectrometer observations of the prototypical starburst galaxy M82. These high resolution spectra represent the best X-ray spectra to date of a starburst galaxy. A complex array of lines from species over a wide range of temperatures is seen, the most prominent being due to Lyman-$\\alpha$ emission from abundant low Z elements such as N, O, Ne, Mg and Si. Emission lines from Helium-like charge states of the same elements are also seen in emission, as are strong lines from the entire Fe$-$L series. Further, the \\Ovii\\ line complex is resolved and is seen to be consistent with gas in collisional ionization equilibrium.  Spectral fitting indicates emission from a large mass of gas with a differential emission measure over a range of temperatures (from $\\sim$0.2\\,keV to $\\sim$1.6\\,keV, peaking at $\\sim$0.7\\,keV), and evidence for super-solar abundances of several elements is indicated. Spatial analysis of the data indicates that low energy emission is more extended to the south and east of the nucleus than to the north and west. Higher energy emission is far more centrally concentrated. ",
        "introduction": "Starbursts play a key role in galaxy formation and evolution, and, in addition, feedback from star-formation in starburst galaxies, in the form of superwinds, probably plays a key role in enriching and heating the intergalactic medium (see Heckman 2001 for a recent review). Further, because superwinds are driven by hot gas generated by stellar winds and supernovae in the starburst, X-ray observations of starbursts have a pivotal position in understanding the structure of starburst and superwinds.  M82 is regarded as one of the archetypical starburst and superwind galaxies. It is a nearby object ($D=3.63$\\,Mpc, Freedman \\etal\\ 1994), is very IR luminous, and is currently undergoing a strong starburst (Rieke \\etal\\ 1993). M82 has been observed by all major X-ray satellites (e.g.\\,Read, Ponman \\& Strickland 1997; Strickland, Ponman \\& Stevens 1997; Cappi \\etal\\ 1999; Kaaret \\etal\\ 2001), and the superwind is clearly visible in X-rays and also in H$\\alpha$ (Shopbell \\& Hawthorn 1998), with clear evidence of a bipolar outflow and a good deal of similarity in the X-ray/H$\\alpha$ morphology (Lehnert, Heckman \\& Weaver 1999). Devine \\& Bally (1999) noted a feature termed the H$\\alpha$ \\lq\\lq cap\\rq\\rq\\ lying well above the galaxy disc that has X-ray emission associated with it.  The {\\em XMM-Newton} satellite has the largest collecting area of any imaging X-ray telescope, and this coupled with the spectral-imaging (EPIC) and grating instruments (Reflection Grating Spectrometers $-$ RGS) offers a new window on starburst galaxies. As M82 is the brightest starburst in the sky, these observations are of key importance to understanding starbursts. In this paper we present results from the RGS. Preliminary results from the EPIC instruments will be presented in Bravo-Guerrero \\etal\\ (2002). It is worth noting that at the present time, almost no RGS results from non-active spiral galaxies have been presented, the most relevant being the very preliminary results from the other famous nearby starburst, NGC~253 (Pietsch \\etal\\ 2001).  The RGS instruments (den Herder \\etal\\ 2001) on board the European X-ray observatory {\\em XMM-Newton} (Jansen \\etal\\ 2001) offer, at least for objects of extent of order 1$-$2\\arcm, the possibility of very high resolution X-ray spectroscopy combined with some spatial resolution in the cross-dispersion direction. Such unique capabilities are excellently suited to the study of the thermodynamic properties of the hot gas within and surrounding the M82 starburst nucleus. Here we describe our preliminary findings in analysing the RGS data obtained during the {\\em XMM-Newton} guaranteed time observations. Results from further more in-depth analysis will be presented in a future paper. The next section describes the observations and the data reduction techniques. Preliminary spectral and spatial results are then presented and discussed, and finally we present our conclusions. ",
        "conclusions": "We have presented here the first high resolution spectra of the M82 starburst taken with the {\\em XMM-Newton} RGS instrument. As M82 is the brightest starburst galaxy and {\\em XMM-Newton} is the satellite with the largest collecting area, these spectra represent the best X-ray spectra to date of a starburst galaxy. Initial spectral and spatial results have been presented and these can be summarized as follows:  \\begin{itemize}  \\item The RGS spectrum is dominated by Ly$\\alpha$ emission lines of hydrogenic charge states of the abundant low Z elements (N, O, Ne, Mg, Si). Helium-like charge states of the same elements are also seen in emission, as is emission from the entire Fe$-$L series, from \\Fexvii\\ through to \\Fexxiv.  \\item The \\Ovii\\ complex is resolved into resonance, intercombination and forbidden lines, and the line ratios are consistent with hot gas in collisional ionization equilibrium. Significant structure is also observed in the \\Neix\\ triple, though contamination from Fe exists.  \\item The \\Fexvii\\ line strengths are consistent with emission from gas in collisional ionization equilibrium. Gas over a range of temperatures (0.3\\,keV to 1.5\\,keV, or higher) can be inferred from the presence of the lines alone.  \\item The M82 starburst appears generally hotter and brighter (by a factor of $\\sim$4) when compared with NGC~253. The weaker longer wavelength lines may suggest a higher level of photoelectric absorption in M82.  \\item A multi-temperature, variable-abundance {\\em mekal} model fits the data well. The differential emission measure shows emission from gas over a range of temperatures (from 0.2\\,keV to 1.6\\,keV, peaking at around 0.7\\,keV). High abundances are obtained for Mg and Si, while near-solar values are obtained for N, O and Fe.  \\item The \\Oviii\\ line profile suggests more \\Oviii\\ emission lying to the east (and south-east) of the nucleus, compared to the west. The more symmetric, thinner higher energy line profiles indicates their emission as being more centrally localised and uniformly distributed.  \\item Cross-dispersion line profiles suggest that emission from the lower energy lines is more extended to the south (and SW). Higher energy lines show similar behaviour, but the inferred extent is reduced.  \\end{itemize}"
    },
    "0206/astro-ph0206464_arXiv.txt": {
        "abstract": "The density irregularities and holes visible in many {\\it Chandra} X-ray images of cluster and galactic cooling flows can be produced by symmetrically heated gas near the central galactic black hole. As the heated gas rises away from the galactic center, a relatively small number of large plumes and bubbles are formed in qualitative agreement with the observed features. The expanding centrally heated gas drives a shock into the surrounding gas, displacing it radially. Both computational and analytic results show that the ambient gas near the bubble is cooled by expansion, accounting for the cool rims commonly observed around X-ray holes in cooling flows. ",
        "introduction": "High resolution X-ray images of ``cooling flows'' in elliptical galaxies taken with the {\\it Chandra} observatory indicate that the central hot gas in these systems is not smoothly distributed, but is cavitated  on scales comparable to the radio emission (M87: B\\\"ohringer et al. 1995; Fabian et al 2001; Wilson et al 2001; M84/NGC 4374: Finoguenov \\& Jones 2001: NGC 4636: Jones et al. 2001; Loewenstein et al. 2001; David et al. 2002; NGC 507: Forman et al. 2001; NGC 5044: Buote et al. 2002). In view of their short dynamical times, $10^7 - 10^8$ yrs, these disturbances, and by inference also the radio sources, must be highly transient. X-ray holes are also seen in galaxy clusters on scales of $\\lta 50$ kpc, often approximately coincident with lobes of extended radio emission (e.g. Hydra A/3C295: McNamara et al. 2000; David et al. 2001; Allen et al 2001; Perseus/NGC 1275/3C84: Churazov et al. 2000; B\\\"ohringer et al. 1993; Fabian et al 2000; A2052: Blanton et al 2001; A4059: Heinz et al 2002). The Perseus cluster in particular contains a multitude of X-ray holes located at random azimuthal orientations. The X-ray holes are regions in which the gas is either heated or displaced by a non-thermal plasma of comparable energy density. In either case the low-density holes must be buoyant (e.g. Churazov et al. 2001).  One of the most remarkable features of these holes is that the X-ray gas around the rims is often {\\it colder} than the average nearby ambient gas (e.g. Fabian et al. 2000, 2001; Fabian 2002). Low temperature rims have been held as evidence that the surrounding gas has not been (recently) strongly shocked, although shocks must accompany even a subsonic expansion of the hot gas. These low entropy rims cannot be understood as local gas that was shocked and subsequently lost entropy by radiation (Nulsen et al 2002; Soker, Blanton \\& Sarazin 2002), but may instead be low-entropy gas that has somehow been raised from the center of the flow as these authors suggest.  The random location of older, more distant bubbles in Perseus and the irregular X-ray images within the central few kpcs of elliptical galaxies. are apparently inconsistent with non-thermal jets having fixed orientations defined by the spin axis of massive black holes. Conversely, not every strong radio lobe corresponds to an X-ray hole, as in Perseus/NGC1275/3C84.  We show here that both the random orientation of X-ray holes and their cold rims can be understood if low entropy hot gas near the center of a cooling flow is heated by an active nucleus in the central elliptical galaxy. ",
        "conclusions": "If central heating is responsible for the frequently observed X-ray holes and surface brightness fluctuations in typical {\\it Chandra} images, as we claim, then such AGN-black hole heating may be a nearly universal component of cooling flows. The physical origin and nature of this heating -- e.g. shock waves, cosmic rays, etc. -- has not yet been explored but these details are not essential to the formation of plumes and bubbles. While the production of X-ray irregularities is quite generally insensitive to the assumed heating parameters, we have noticed several trends when $L_h$, $r_h$ and $t_h$ are varied. Non-linear plumes mature faster when $r_h$ is reduced. If $r_h$ is too large, the cold rims are less pronounced. However, the possibility of learning about such heating details by comparing X-ray images with gasdynamical calculations is limited if the mean time between heating episodes is sufficiently short that the initial pre-heated gas is no longer perfectly quiescent as we have supposed here.  All well-observed cooling flows emit diffuse optical Balmer and forbidden lines visible in the brightest central regions. Enhanced H$\\alpha$ + [NII] emission has been observed near the boundaries of X-ray bubbles (e.g. Blanton et al 2001; Trinchieri \\& Goudfrooij 2002). It is tempting to conclude that such gas has cooled from the hot phase in the cooler, denser gas around the rims of the plumes or bubbles. But direct cooling from the hot phase is not supported by our detailed hydrodynamic calculations which are essentially unchanged if radiative cooling losses are neglected. The radiative cooling time is greater than the total time of our simulations. The cooler gas at $T \\sim 10^4$ K responsible for the observed optical line emission may simply be displaced by the plumes and bubbles along with the local hot gas. However, on longer timescales than we consider here, heating in the central regions stimulates cooling by thermal instabilities in more distant regions of the flow, so the time-averaged global cooling rate is essentially unchanged (Brighenti \\& Mathews 2002).  If intermittent AGN heating is reasonably frequent, the entire inner cooling flow will become convective and turbulent (Brighenti \\& Mathews 2002). Turbulent motions observed in the diffuse optical lines (e.g. Caon et al. 2000) are too energetic on large scales to be understood in terms of supernova explosions, but could result from central heating similar to that described here.  X-ray holes formed by central heating are randomly disposed around the center of the flow, as in the Perseus cluster. This is a particularly desirable feature since the X-ray bubbles need not be aligned with the accretion axis of the central black hole. Nevertheless, radio lobes often appear to be associated with X-ray holes. This connection may be causal, but in some cases the relativistic electrons may readily flow into pre-existing holes; this may be occurring in the central southwestern hole in Perseus (Fig. 1 of Fabian 2001).  Finally, the azimuthally averaged hot gas density profile in these centrally heated models is flattened by the holes, providing an excellent fit to observed X-ray surface brightness profiles which are too centrally peaked otherwise. To maintain this fit, AGN heating would need to reoccur every $\\sim 10^8$ years. However, in the presence of central heating the gas temperature profile is less satisfactory, lacking the deep central minimum usually observed.  We conclude that central heating in cooling flows naturally results in a relatively small number of randomly located X-ray bubbles similar to the irregularities observed. In addition, these expanding plumes and bubbles are surrounded by considerably cooler gas, also consistent with recent {\\it Chandra} observations. Gas in the cold rims was shock-heated but subsequently cooled by expansion.   \\vskip.4in Studies of the evolution of hot gas in elliptical galaxies at UC Santa Cruz are supported by NASA grant NAG 5-8049 and NSF grants AST-9802994 and AST-0098351 for which we are very grateful. FB is supported in part by grants MURST-Cofin 00 and ASI-ARS99-74."
    },
    "0206/astro-ph0206008_arXiv.txt": {
        "abstract": "We investigate the spatial, kinematic and chemical properties of globular cluster systems formed in merging and interacting galaxies using N-body/SPH simulations. Although we can not resolve individual clusters in our simulation, we assume that they form in collapsing molecular clouds when the local external gas pressure exceeds 10$^5$ $k_B$ (where $k_B$ is the Boltzmann constant). Several simulations are carried out for a range of initial conditions and galaxy mass ratios. The input model spirals are given a halo globular cluster system similar to those observed for the Milky Way and M31. Gravitational tidal effects during galaxy merging and interaction leads to a dramatic increase in gas pressure, which exceeds our threshold and hence triggers new globular cluster formation. We investigate the properties of the globular cluster system in the remnant galaxy, such as number density, specific frequency, kinematic properties and metallicity distribution. Different orbital conditions and mass ratios give rise to a range in globular cluster properties, particularly for the interaction models. Our key results are the following: The newly formed metal-rich clusters are concentrated at the centre of the merger remnant elliptical, whereas the metal-poor ones are distributed to the outer parts due to strong angular momentum transfer. The dissipative merging of {\\it present day} spirals, including chemical evolution, results in metal-rich clusters with a mean metallicity that is super-solar, i.e. much higher than is observed in elliptical galaxies. If elliptical galaxies form by dissipative major mergers, then they must do so at very early epochs when their discs contained low metallicity gas. Our simulations show that the specific frequency can be increased in a dissipative major merger. However, when this occurs it results in a ratio of metal-poor to metal-rich clusters is less than one, contrary to the ratio observed in many elliptical galaxies. ",
        "introduction": "There have been many previous models and scenarios for the formation of globular clusters and their systems (e.g., Peebles \\& Dicke 1968; Searle \\& Zinn 1978; Fall \\& Rees 1985; Fall \\& Rees 1988; Zinnecker, Keable, \\& Dunlop 1988; Larson 1987, 1988; Kang et al. 1990; Ashman \\& Zepf 1992; Freeman 1993; Harris \\& Pudritz 1994; Elmegreen \\& Efremov 1997; Forbes et al. 1997; McLaughlin 1999; Ashman \\& Zepf 2001; Bekki \\& Couch 2001; Cen 2001; Weil \\& Pudritz 2001; Beasley et al. 2002 and Bekki \\& Chiba 2002: See Harris 1991 and Ashman \\& Zepf 1998 for a review). These studies have contributed to a  better understanding of the physical conditions of globular cluster formation at low and high redshifts, in addition to providing insight into the the physical origins of the observed scaling relations of globular cluster systems.  However, they have not addressed the observed structural, kinematical, and chemical properties of globular cluster systems {\\it in a self-consistent manner} through the detailed modelling of collapsing  molecular clouds. Furthermore, given the observational fact that young super star clusters are often located in star-burst regions (e.g., Ashman \\& Zepf 1998; 2001), it is critically important to investigate physical properties of globular clusters formed by the induced collapse of molecular clouds {\\it for the case of galaxy merging and interaction}. Comparing the predicted properties of globular cluster systems with the corresponding observational ones in E/S0 galaxies is important, because recent  observations provide a  rich data set of metallicity and kinematics for  globular  clusters around giant elliptical galaxies (e.g., Kissler-Patig et al. 1999; Cohen et al. 1998 Beasley et al. 2000; Bridges 2001; Forbes 2001; Forbes et al.  2001). Furthermore  better understanding   of the nature  of globular cluster formation in {\\it  present-day} merging and interacting galaxies can provide clues to their formation {\\it at  high redshift}, where merging and interaction was  much more frequent.     The purpose of this paper is to numerically investigate the kinematical and chemical properties of globular clusters formed in merging and interacting spiral galaxies. We adopt the plausible assumption that the high pressure of warm interstellar gas ($P_{\\rm gas}$ $>$ $10^5$ $k_{\\rm B}$: $k_{\\rm B}$ is Boltzmann's constant) can induce the global collapse of giant molecular clouds to form massive compact star clusters corresponding to super star clusters or progenitor objects of globular clusters (Jog \\& Solomon 1992; Elmegreen \\& Efremov 1997).  We mainly investigate the following five points: (1) why is globular cluster formation more efficient in merging galaxies than in isolated spiral galaxies, (2) how does the specific frequency ($S_{\\rm N}$) of globular cluster systems change during merging, (3) what are the fundamental properties (e.g., structure, kinematics, and metallicity distribution) of newly formed globular cluster systems, (4) how do the physical properties of globular clusters formed during merging depend on orbital configuration, mass-ratio of the merging spirals, gas mass fraction, and merging epoch, and (5) are there any differences in the details of globular cluster formation in tidally interacting galaxies versus merging galaxies. We also stress that better understanding the effects of galactic tides on the dynamical and chemical evolution of interstellar medium is of  primary importance  for clarifying  the unresolved problems related to globular cluster formation in the low and high redshift universe.     The plan of the paper is as follows: In the next section, we describe our  numerical model for  globular cluster formation based on the adopted molecular cloud collapse scenario. In \\S 3, we present the numerical results on structural, kinematical, and chemical properties of globular clusters formed in merging and interacting  galaxies. In \\S 4, we  compare the basic properties of globular cluster systems from our model predictions with the observations. We also present additional predictions which can be tested against future observations. We summarise our  conclusions in \\S 5. ",
        "conclusions": "We have numerically investigated how the global dynamical evolution of merging and interacting galaxies can determine the formation processes and fundamental physical properties of globular clusters.  We have assumed that a globular cluster can be formed from a giant molecular cloud in a galaxy if the warm interstellar gas of the galaxy becomes high enough ($P_{\\rm gas}$ $>$ $10^5$ $k_{\\rm B}$).  We summarise our principle results as follows.   (1) Strong tidal shocks induced by galaxy merging and interaction can dramatically increase gaseous pressure ($P_{\\rm g}$ $>$ $10^5$ $k_{\\rm B}$) so that molecular clouds can collapse to form globular clusters. During the formation of globular clusters in a merging/interacting galaxy, the ratio of the formation rate of globular clusters to that of field stars increases due to the larger fraction of gas with high pressure in the galaxy. Thus globular cluster formation is more efficient in star-burst regions of galaxies. This result can explain why young globular clusters are commonly observed in merging and interacting galaxies compared to isolated spirals.   (2) The specific frequency ($S_{\\rm N}$) of globular cluster systems can be increased a factor of $2-3$ in a major gaseous merger (which results in the formation of an elliptical galaxy) due to the creation of new metal-rich globular clusters. However, many elliptical galaxies are observed to have higher $S_{\\rm N}$ values and higher ratios of metal-poor to metal-rich clusters than can be explained by our merger simulations. Merger progenitor spirals with $S_{\\rm N}$ values higher than those assumed here (i.e., 0.8) may help aleviate this problem.      (3) The metallicity distribution of metal-rich globular cluster systems formed by major merging depends upon the initial metallicity distribution of merger progenitor discs and the chemical evolution of gas. In present-day mergers, the mean value of the metallicity distribution of newly formed globular clusters is typically higher than twice solar metallicity because of the initially large gas metallicity and efficient chemical processing. This means that recent major mergers can yield bimodal globular cluster colour distributions but, these distributions have metal-rich mean values that are much greater than that observed in elliptical galaxies(i.e, [Fe/H] $\\sim$ $-0.5$). This suggests that if most of elliptical galaxies are formed by major merging, then the typical merging epoch should be a high redshift when the merger progenitor discs have low-metallicity gas.   (4) The dynamical evolution of metal-poor globular clusters (initially in the halo regions of merger progenitor spirals) is different from that of the newly formed metal-rich clusters. Metal-poor globular clusters experience stronger angular momentum transfer from the inner to outer regions, whereas metal-rich clusters experience a larger amount of gas dissipation prior to their formation.  As a natural result of this, the surface density distribution and kinematical properties of the two globular cluster subpopulations are different. For example, the metal-poor globular clusters show shallower density profiles, larger velocity dispersion in the central regions and a large amount of rotation in the outer regions.   (5) The $S_{\\rm N}$, structural, and kinematical properties of globular clusters in the remnants of major mergers depend weakly on the orbital configurations of the merger (such as initial orbital angular momentum and whether a merger is prograde-prograde or retrograde-retrograde).  The metallicity distribution of the newly formed globular cluster system does not depend on the above orbital parameters. Given the observed positive correlation between luminosity and gas metallicity in disc galaxies (Zaritsky et al. 1994), our results suggest that there should be a positive correlation between metallicities of metal-rich globular clusters and luminosities of their host galaxies.   \\indent (6) The formation efficiency, total number, and $S_{\\rm N}$ of globular clusters formed in mergers depends on the mass ratio of the two merging discs ($m_{2}$) in such a way that each of these three quantities are smaller in mergers with smaller $m_{2}$. This is due to the weaker galactic tide in mergers with smaller $m_{2}$, as only a smaller fraction of gas can have enough high gas pressure ($P_{\\rm g}$ $>$ $10^5$ $k_{\\rm B}$) to trigger the molecular cloud collapse leading to globular cluster formation. These results imply that S0s, which can be formed by minor and unequal-mass merging, will show $S_{\\rm N}$ values that are appreciably higher than that of typical spirals, but lower than that of ellipticals.  These results further suggest that spiral galaxies with thick disc components, which can be formed by minor merging, should also show $S_{\\rm N}$ values slightly higher than that of spirals without thick discs.\\\\  \\indent (7) The $S_{\\rm N}$ of globular clusters is more likely to be lower in tidally interacting galaxies than in merging ones.  The formation efficiency, total number, and $S_{\\rm N}$ of globular clusters in interacting galaxies strongly depends on the structure of disc galaxies, orbital configurations, and the mass ratio of two interacting discs.  For example, $S_{\\rm N}$ does not increase significantly in low surface-brightness galaxies compared to high surface-brightness galaxies. Furthermore $S_{\\rm N}$ is higher in an interacting galaxy with a larger mass-ratio. As tidal interaction may also transform gas-rich spirals into a gas-poor S0s, our tidal interaction models can provide a natural explanation for the observed $S_{\\rm N}$ values of S0s.\\\\  \\indent (8) The formation efficiency of globular clusters in merging and interacting galaxies is likely to decrease with galaxy mass, though field star formation is still ongoing in those galaxies.  This implies that there could be a minimum galactic mass for a galaxy to harbour globular clusters.  This result also provides an important implication for the formation of metal-poor halo globular clusters that are believed to be formed more than 10 Gyr ago, if such old globular clusters are formed from merging of less massive sub-galactic gaseous clumps at high redshift."
    },
    "0206/astro-ph0206244_arXiv.txt": {
        "abstract": " ",
        "introduction": "``Galaxies: the Third Dimension'' continues a tradition which started in Marseilles 1994 with the ``Tridimensional Optical Spectroscopic Methods in Astrophysics'' meeting, followed by ``Imaging the Universe in Three Dimensions'' at Walnut Creek in 1999.  We were both fortunate to attend all three meetings, and this review provides us with the opportunity for retrospect. There have been developments both in instrumentation and in the way astronomy is conducted, and these need to be seen against a broad canvas. As we emphasize below, one of the main reasons for satisfaction is the trend towards talks focused on scientific results rather than on `yet another 3D spectroscopic technology.' ",
        "conclusions": ""
    },
    "0206/astro-ph0206072_arXiv.txt": {
        "abstract": "Soft Gamma Repeaters undergo pulse profile changes in connection with their burst activity. Here we present a comprehensive pulse profile history of SGR 1806-20 and SGR 1900+14 in three energy bands using Rossi X-ray Timing Explorer/Proportional Counter Array observations performed between 1996 and 2001. Using the Fourier harmonic powers of pulse profiles, we quantify the pulse shape evolution. Moreover, we determined the RMS pulsed count rates (PCRs) of each profile. We show that the pulse profiles of SGR 1806$-$20 remain single pulsed showing only modest changes for most of our observing span, while those of SGR 1900+14 change remarkably in all energy bands. Highly significant pulsations from SGR 1900+14 following the 1998 August 27 and 2001 April 18 bursts enabled us to study not only the decay of PCRs in different energy bands but also their correlations with each other. ",
        "introduction": "Soft gamma repeaters (SGR) constitute a small class of isolated neutron stars; four are currently well-established while one more source remains an unconfirmed candidate (Cline et al. 2000). As their name indicates, SGR emission is repetitive: during their active states each source emits hundreds of short (duration $\\sim$ 0.1 s), intense (at $\\sim$ super-Eddington luminosities) bursts of hard X-rays / soft gamma-rays at random intervals. Very rarely, (observed once so far from each of SGR~$0526-66$ and ~$1900+14$) they emit a giant flare, an event with $\\sim$10$^4$ higher luminosity than the typical short events (Mazets et al. 1979, Hurley et al. 1999a, Feroci, et al. 2001). Thompson \\& Duncan (1995; henceforth TD95) proposed that the SGR burst activity is associated with the strong magnetic field (B$\\gtrsim$ 10$^{14}$ G) of the neutron star. Their model, also known as the magnetar model, attributes SGR bursts to various levels of fracturing of the neutron star crust by the motion of the anchored lines of their strong magnetic fields (TD95, Thompson \\& Duncan 2001).  SGR sources were established as a new class in the mid 80s based on their bursting properties. The first X-ray counterpart of an SGR was discovered in 1993 with the Japanese satellite ASCA (Murakami et al. 1994) after SGR~$1806-20$ triggered the Burst And Transient Source Experiment (BATSE) onboard the Compton Gamma Ray Observatory (CGRO) (Kouveliotou et al. 1994). All four SGRs are currently associated with persistent X-ray point sources with luminosities 10$^{34}$ $-$ 10$^{35}$ ergs s$^{-1}$. Their energy spectra are nominally fit with a single power law with photon indices $2-3.2$, except in the case of SGR~$1900+14$, where a two component spectrum (blackbody of $kT=0.50(4)$ and power law with index $=2.1(3)$) has been established with BeppoSAX and Chandra observations (Woods et al. 1999a, Kouveliotou et al. 2001).  Observations of SGR~$1806-20$ with the Rossi X-Ray Timing Explorer (RXTE) Proportional Counter Array (PCA) on November 1996 led to the discovery of the first SGR spin period of 7.47 s. Further RXTE/PCA observations established that the source exhibited a very rapid spin-down rate of 8.1 $\\times$ 10$^{-11}$ s s$^{-1}$, providing the first direct measurement of the neutron star magnetic field of $\\sim$ 2$\\times$ 10$^{14}$ G\\footnotemark, \\footnotetext{Assuming that the star slows down via magnetic dipole radiation, B$_{\\rm d}$ $\\propto$ $\\sqrt{{\\rm P}\\dot{\\rm P}}$.} in agreement with the magnetar model of TD95 (Kouveliotou et al.~1998a). ASCA observations of SGR~$1900+14$ in 1998 led to the discovery of a spin period for that source (Hurley et al. 1999b) and to the detection of its rapid ($1.1\\times10^{-10}$ s s$^{-1}$) spin down rate (Kouveliotou et al. 1999). Of the two remaining SGRs, SGR~$0526-66$ has been associated with an 8 s spin period, originally seen as intensity modulation in the decaying tail of the giant flare it emitted in 5 March 1979 (Mazets et al. 1979). A pair of recent Chandra observations of this source during quiescence spaced $\\sim$ 1 year apart should establish the periodicity and spin down rate of SGR~$0526-66$ (Kulkarni et al. 2002). The last source, SGR~$1627-41$, is the dimmest of them all and no spin modulation has been detected in its persistent X-ray flux (Kouveliotou et al. 1998b).  In 1998 we started a monitoring campaign with RXTE/PCA for the two SGR sources with clearly established spin and spin down rates. Our observations so far have shown that these spin down rates vary significantly, sometimes by a factor of four, and that they exhibit a very high level of timing noise (Woods et al. 2000, 2002). Earlier results have also indicated a relationship between the pulse profile complexity and the activity history of the source (Woods et al. 2001).  In this paper we construct the first detailed history of the pulse profile evolution of SGR~$1806-20$ and SGR~$1900+14$ spanning the last $\\sim$ 5 years, and we examine their evolution with energy. We quantify the profile changes by estimating the power in their respective Fourier harmonics. In Section 2, we describe our observations and how we deal with the X-ray background component in our data. In Section 3, we describe the methodology of our data analysis and present the results; we discuss the implications of our results in Section 4. ",
        "conclusions": "We have performed a detailed analysis of the persistent X-ray pulsed flux data from two Soft Gamma Repeaters, SGR 1900+14 and SGR 1806-20. In both sources, we find a strong trend of their pulse shapes to transition from complex to very smooth, following burst active episodes. Our analysis of the decaying tail of the 27 August giant flare from SGR 1900+14, has shown that as the persistent source spectrum softened, the pulse count rates in the 2-5 keV band varied in tandem with those of the 5-10 keV; the 10-20 keV pulsed count rates remained constant, albeit slightly above their pre-flare values. Further, we observe a change in the correlation slope of the pulsed count rates in the lower energy ranges (2-5, and 5-10 keV) during quiescence or low activity of the source, suggesting that the 2-5 keV photons are more efficiently produced during non-burst active periods.  Recently, Thompson, Lyutikov \\& Kulkarni (2002; TLK hereafter) have suggested a model, whereby the persistent X-ray emission of a magnetar is due to currents generated when the interior magnetic field twists up the external magnetic field. The structure we observe in the pulse profile of SGRs could then be attributed to a complicated distribution of currents and magnetic fields close to the neutron star (higher multipoles) seen through a largely transparent magnetosphere. However, the observed transition of complex to smooth in the pulse profiles is not accompanied by spectral and intensity changes, as would be expected in that case from a simplified magnetic field (current) after the source returns to its quiescent flux level. Alternatively, re-scattering of the X-rays either at $R>100$ km by non-axisymmetric currents, or at ~100 km by plasma suspended against gravity by the resonant e$^{-}$ scattering force, may be the cause of the profile smoothness. The apparent difficulty to retain the suspended plasma in the magnetosphere (it would be quickly drained to the neutron star surface by even a modest electrical current), points towards the re-scattering screen above 100 km as the most plausible of these two explanations for the transition. The current then would in principle be generated by a static twist (TLK 2002), consistent with the observed absence of a direct correlation between the increase in the neutron star spin (torque) with burst activity (Woods et al. 2002).  Concluding, we believe that our results provide evidence for an association of the complex(smooth) profile of the SGR pulses with the re-scattering of X-rays above 100 km from the neutron star surface, depending on whether the currents created by the static twist described by TLK (2002) are non-axisymmetric(axisymmetric). The rapid spindown observed in SGRs should then be coupled to the smoothening of their pulse profiles; we are currently investigating further correlations between all these SGR properties."
    },
    "0206/astro-ph0206302_arXiv.txt": {
        "abstract": "The SN explosion in the closed binary can give the magnetospheric flare possessing the properties of GRB. The SN shock, flowing around the magnetosphere of a magnetized neutron star or a white dwarf, produces a narrow magnetic tail $10^9 cm$ long, $10^8 cm$ wide and a magnetic field of $10^6 Gauss$. Fast particles ( $\\gamma\\simeq 10^4$ ), generated in the tail by reconnection processes, radiate gamma rays of the $100KeV$ - $1MeV$ energies. The duration of radiation $T<1 sec$ corresponds to a short GRB. Apart, the powerful shock can tear and accelerate part of the tail. That is the relativistic ( $\\Gamma\\simeq 10^4$ ), strongly magnetized jet, producing gamma radiation and also X-ray and optic afterglow. That is a long ($T>10 sec$) GRB. The duration of the afterglow is inversely proportional to the photon energy and is several months for optic. ",
        "introduction": "Identification of two tens of the long GRB ( $T_{90}>1sec$ ) with the extragalactic objects located at cosmological distances ( $z=0.1 - 4$ ) raises a question about the source of energy release. For the spherical symmetry the energy yield is turned to be as abnormally large as $10^{53-54} ergs$ only in the band of thousands of $KeV$. Such energy release is comparable with that of a supernova explosion ( SN ), for which $90\\%$ of the energy is escaped by neutrinos during the first several seconds. To overcome the 'energy catastrophe' in the problem of GRB a hypothesis was suggested about a strong collimation of the gamma radiation in the narrow solid angle of $\\Delta\\phi\\simeq 1^\\circ$. This assumption gives a possibility to diminish the luminosity of GRB, needed for the explanation of the observed gamma-ray flux, by a factor of $ (\\Delta\\phi/4)^2 $ which for $\\Delta\\phi\\simeq 1^\\circ$ gives the gain of $10^5$. Such directivity can be explained by the radiation of a narrow jet of particles accelerated to relativistic energies. However, it is not clear now how this jet is formed when two compact objects merge or when one of them explodes. That is why there exists a large amount of models suggested for GRB explanation, \\cite{Djo}. In last years the model of strong anisotropic supernova explosion becomes very popular ( see, for example, \\cite{Pac} ), which demands extreme conditions for its realization.  Our model suggests the mechanism of origin of a narrow beam of gamma-rays when a shock of an SN explosion in the binary system, containing a magnetized neutron star or a white dwarf, is interacting with its magnetosphere. If the observer is placed in the plane of the orbit of such a binary system and at the moment of the origin of the magnetospheric tail the line of observation coincides with the tail direction, then a GRB will be observed. We give the estimations confirming that our model can explain the general properties of GRB, both long ( $T>1sec$ ) -- cosmological, and short ( $T<1sec$ ) -- possibly local. For the short GRB we consider the binary system in which one companion, as previously, is a neutron star or a white dwarf and another is a star of cataclysmic flare variable. ",
        "conclusions": ""
    },
    "0206/astro-ph0206134_arXiv.txt": {
        "abstract": "{ A nonlinear nature of the binary ephemeris of Cygnus~X-3 indicates either a change in the orbital period or an apsidal motion of the orbit. We have made extended observations of Cygnus~X-3 with the Pointed Proportional Counters (PPCs) of the Indian X-ray Astronomy Experiment (IXAE) during 1999 July 3$-$13 and October 11$-$14. Using the data from these observations and the archival data from ROSAT, ASCA, BeppoSAX and RXTE, we have extended the data base for this source. Adding these new arrival time measurements to the published results, we make a comparison between the various possibilities, (a) orbital decay due to mass loss from the system, (b) mass transfer between the stars, and (c) apsidal motion of the orbit due to gravitational interaction between the two components. Orbital decay due to mass loss from the companion star seems to be the most probable scenario. ",
        "introduction": "The nature of the compact object in the bright X-ray binary Cygnus~X-3, is a subject of much debate. In spite of being one of the most frequently observed X-ray sources, presence of an X-ray pulsar, a black hole, or a low magnetic field neutron star has not yet been established. There is also uncertainty about the mass and type of the companion star. An interesting way to probe this system is to investigate the arrival time history of the 4.8 hr orbital modulation in the X-ray light curve.  The unusual X-ray binary Cygnus~X-3 is located in the plane of our galaxy at a distance of $>$ 11.6 kpc (Dickey 1983). Because of the high luminosity of the source in X-ray, infrared and radio bands, it has been observed on many occasions at these wavelengths. Strong optical extinction in the direction of the source prevents optical observations. The mass of the companion star has been determined by many different means  and it is found to be in the range of a fraction of solar mass to a few solar mass (Van den Heuvel \\& De Loore 1973; Tavani, Ruderman, \\& Shaham 1989; van Kerkwijk et al. 1992). Parsignault et al. (1976) discovered a periodicity of 4.8 hr in the X-ray flux with nearly sinusoidal variation, which is believed to be due to the orbital motion. The variation in the sinusoidal shape of the light curve from cycle to cycle is reported by van der Klis and Bonnet-Bidaud (1981). They showed that the shape of the average light curve formed from successive cycles is quasi-sinusoidal with a slow rise and fast fall. Several models have been proposed to explain the deep and near-sinusoidal modulation of the X-ray intensity of Cygnus~X-3 with the orbital period (Ghosh et al. 1981). Material in the form of a large shell around the binary system can produce the X-ray modulation by scattering, absorbing, and re-emitting the X-rays from the compact object, if the X-ray emission is only from the non-shadowed part of the shell. White \\& Holt (1982) proposed the accretion disk corona (ADC) model to explain the modulation in the X-ray light curve of the source. According to this model, the X-ray source is covered by an optically thick corona with a radius of about 10$^9$ cm, an optical depth of about 10 and a temperature of about 2 keV. Another model which tries to explain the orbital modulation is stellar wind model (Willingale, King, \\& Pounds 1985; Kitamoto et al. 1987). This model assumes a strong and highly ionized stellar wind from the companion star with an optical depth of about 1 and a radius of about 10$^{11}$ cm and the temperature of the wind is model dependent.  Evolution of the orbit of Cygnus~X-3 is studied by measuring the arrival times of the minima in each orbital motion. The ephemeris of 4.8 hr modulation in the X-ray light curve of the source has been studied by many authors (Leach et al. 1975; Mason \\& Sanford 1979; Parsignault et al. 1976; Lamb, Dower \\& Fickle 1979; Elsner et al. 1980; van der Klis \\& Bonnet-Bidaud 1981, 1989; Kitamoto et al. 1987). The time derivative of the orbital modulation period ($\\dot{P}$) and second derivative ($\\ddot{P}$) have been measured as $\\sim$ 10$^{-9}$ s s$^{-1}$ and $\\sim$ -10$^{-11}$ yr (van der Klis \\& Bonnet-Bidaud 1981, 1989; Kitamoto et al. 1987). The unusually large value of $\\dot{P}$/P (= 2.2 $\\times$ 10$^{-6}$ yr$^{-1}$ ), which if linked to the binary evolution of the system, is similar to several short-period bright LMXBs such as X1822$-$371 with $\\dot{P}$/P $\\sim$ 3.4 $\\times$ 10$^{-7}$ yr$^{-1}$ and orbital period of 5.57 hr (Hellier et al. 1990; Parmar et al. 2000), EX00748$-$767 with $\\dot{P}$/P $\\sim$ 2 $\\times$ 10$^{-7}$ yr$^{-1}$ and orbital period of 3.82 hr (Parmar et al. 1991). This value of $\\dot{P}$ can be interpreted as due to the mass transfer from the optical companion to the compact object. Using the measured period derivative of Cygnus X$-$3, Kitamoto et al. (1987) determined the rate of mass loss from the binary companion as about 10$^{-6}$ M$_{\\odot}$ yr$^{-1}$. Tavani et al. (1989) tried to explain the observed high value of $\\dot{P}$/P in Cygnus~X-3 and the required mass transfer rate close to $\\sim$ 1.58 $\\times$ 10$^{-8}$ M$_{\\odot}$ yr$^{-1}$. They claimed that the companion is a degenerate star with solar composition and in the mass range of 0.01 $\\leq$ $m$ $\\leq$ 0.03 M$_{\\odot}$ which under-fills its Roche lobe. The X-ray illumination from the primary compact object produces the evaporative wind from such a low-mass degenerate companion. However, from infrared observations, van Kerkwijk et al. (1992) discovered that the binary companion of Cygnus~X-3 has the spectrum of a Wolf-Rayet star and predicted the mass of the companion as $\\sim$ 10 M$_{\\odot}$ which suggests that Cygnus~X-3 is a high mass X-ray binary. Recent observations by Fender et al. (1999) confirm the nature of the companion as an early-type WN Wolf-Rayet star.  We have measured the available arrival time of minimum value of the light curve from two observation of Cygnus~X-3 with PPCs of IXAE in 1999 and several archival data. By combing these new measurements with the previously published results, we are able to determine the rate of change of orbital period with greater accuracy. ",
        "conclusions": "\\begin{itemize} \\item We have measured 34 new minima and extended the measurement base to 30 years.  \\item The parabolic model provides the best fit to the new data with a modified value of $\\dot{P}$.  \\item Sinusoidal model and the corresponding apsidal motion scenario is ruled out.  \\item Mass loss from a massive companion is most favoured in Cygnus~X$-$3. \\end{itemize}"
    },
    "0206/astro-ph0206120_arXiv.txt": {
        "abstract": "{ We consider a sample of 22 nearby clusters of galaxies observed with the Medium Energy Concentrator Spectrometer (MECS) on board \\sax. They cover the range in gas temperature between 3 and 10 keV, with bolometric X-ray luminosity between $2 \\times 10^{44}$ erg s$^{-1}$ and $6 \\times 10^{45}$ erg s$^{-1}$. Using the de-projected gas temperature and density profiles resolved in a number of bins between 5 and 7 and obtained from this dataset only, we recover the total gravitating mass profiles for 20 objects just applying the (i) spherical symmetry and (ii) hydrostatic equilibrium assumptions. We investigate the correlations between total mass, gas temperature and luminosity at several overdensities values and find that the slopes of these relations are independent of the considered overdensity and consistent with what is predicted from the cluster scaling laws. The best-fit results on the normalization of the $M-T$ relation are slightly lower, but still consistent considering the large errors that we measure, with hydrodynamical simulations. A segregation between relaxed and non-relaxed systems is present in each plane of these relations pointing out a significant component in their intrinsic scatter. This segregation becomes more evident at higher overdensities and when physical quantities, like $M_{\\rm gas}$ and $L$, that are direct functions of the amount of gas observed, are considered. ",
        "introduction": "The amplitude and the shape of the power spectrum of the primordial density fluctuations on scales of about 20 $h_{50}^{-1}$ Mpc can be effectively constrained with the mass function of galaxy clusters. Since the early '90s, X-ray observations have been used to build large datasets of measured luminosities and, with more effort because a larger number of source counts is required, temperatures of the X-ray emitting plasma trapped in the cluster gravitational potential. These observed quantities are expressions of the physical processes that are taking place in the galaxy clusters and manifest the energy and the mass of these systems. Then, comparing the observed distribution in luminosity (or temperature) with theoretical models of the expected cluster number density that are functions of total mass and redshift and depend upon the cosmological model adopted (e.g. Press \\& Schechter 1974), it has been possible to put constraints in the ``normalization--shape'' plane of the primordial density fluctuations spectrum (see, e.g., the pioneering work of Henry \\& Arnaud 1991 and the most recent results in Ikebe et al. 2002 and references therein).  However, the conclusions reached making this comparison rely on an efficient way to relate the observed quantities (like gas luminosity and temperature) to the gravitating mass of the systems. Gas--dynamics simulations (e.g. Evrard, Metzler \\& Navarro 1996, Schindler 1996) have confirmed the expected correlation between mass and temperature and have shown that mass estimates are reliable when obtained through X-ray analysis under the assumption of spherical symmetry and hydrostatic equilibrium. More recently, mass profiles obtained relaxing the condition of plasma isothermality have shown a significant mismatch in normalization and slope of the mass--temperature relation between observational data and simulations (e.g. Horner, Mushotzky \\& Scharf  1999, Nevalainen, Markevitch \\& Forman 2000). On the other hand, it has been clear since the first compilation of catalogues of luminosity and temperature (Mushotzky 1984, Edge \\& Stewart 1991) that the observed correlation between these two quantities deviates significantly from the expected scaling law, suggesting contributions to the total energy  of the plasma from  physical phenomena other than the gravitational collapse.  The observed Luminosity-Temperature ($L-T$) and Mass-Temperature ($M-T$) relations for galaxy clusters are, therefore, the foundation to construct the cluster mass function and to use these virialized objects as cosmological probes. In this paper, we investigate these relations and, more in general, any correlation between observed and inferred quantities using \\sax observations of 22 nearby clusters of galaxies with resolved gas temperature and density profiles. The main differences between this study and previous work on the same subject are: \\begin{enumerate} \\item the use of \\sax data that allows us to extend the analysis of spatially-resolved spectra up to 20\\arcmin\\ in radius, i.e. $\\sim 2.5$ times the most favourable configuration with \\chandra (Weisskopf et al. 2000) and to put under control some systematic effects (e.g., sharper and more energy-independent Point-Spread-Function than the \\asca one --Tanaka et al. 1994--, more stable and lower background than the one observed in \\xmm\\ --Jansen et al. 2001), \\item the direct deprojection of the spectral results to reconstruct the gas temperature and density profiles in a model-independent way. \\end{enumerate}  The sample presented in this work is, to date, the largest for which the physical quantities (i.e. gas density, temperature, luminosity, total mass, etc.) have all been derived simultaneously from spatially-resolved spectroscopy of the same dataset. The difference between this approach and  others which make use of data coming from different satellites and/or make strong assumptions on the temperature profiles, such as isothermality, is twofold: on one side the use of data from different missions and the simplistic assumptions on the temperature profiles allow to build up samples bigger than ours, on the other they increase the likelihood of systematic effects which may in turn affect the relations between the observed quantities.  The paper is organized as follows: in Section~2, we describe the \\sax MECS observations of the galaxy clusters in our sample and the results of the spectral analysis considered in this work; the deprojection technique applied to the projected spectral results is discussed in Sect.~3; the gravitating mass profiles are obtained and compared with the optical measurements in Section~4; in Section~5, we study the correlation between the total mass, gas temperature, gas mass and luminosity; we summarize our results and present our conclusions in Sect.~6.  All the errors quoted are at $1 \\sigma$ level (68.3 per cent level of confidence for one interesting parameter) unless otherwise stated. The cosmological parameters $H_0 = 50 h^{-1}_{50}$ km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_{\\rm m} = 1 -\\Omega_{\\Lambda} = 1$ are assumed hereafter. ",
        "conclusions": ""
    },
    "0206/astro-ph0206316_arXiv.txt": {
        "abstract": "{ We present the analysis of multi-wavelength XMM-Newton data from the Seyfert galaxy NGC 3783, including UV imaging, X-ray and UV lightcurves, the 0.2$-$10~keV X-ray continuum, the iron K$\\alpha$ emission line, and high-resolution spectroscopy and modelling of the soft X-ray warm absorber. The 0.2$-$10~keV spectral continuum can be well reproduced by a power-law at higher energies; we detect a prominent Fe K$\\alpha$ emission line, with both broad and narrow components, and a weaker emission line at 6.9 keV which is probably a combination of Fe K$\\beta$ and \\ion{Fe}{xxvi}.  We interpret the significant deficit of counts in the soft X-ray region as being due to absorption by ionised gas in the line of sight. This is demonstrated by the large number of narrow absorption lines in the RGS spectrum from iron, oxygen, nitrogen, carbon, neon, argon, magnesium, silicon and sulphur. The wide range of iron states present in the spectrum enables us to deduce the ionisation structure of the absorbing medium. We find that our spectrum contains evidence of absorption by at least two phases of gas: a hotter phase containing plasma with a log ionisation parameter $\\xi$ (where $\\xi$ is in erg cm s$^{\\rm -1}$) of 2.4 and greater, and a cooler phase with log $\\xi$ centred around 0.3. The gas in both phases is outflowing at speeds of around 800 km s$^{\\rm -1}$. The main spectral signature of the cold phase is the Unresolved Transition Array (UTA) of M-shell iron, which is the deepest yet observed; its depth requires either that the abundance of iron, in the cold phase, is several times that of oxygen, with respect to solar abundances, or that the absorption lines associated with this phase are highly saturated. The cold phase is associated with ionisation states that would also absorb in the UV. ",
        "introduction": "NGC 3783 is a Seyfert 1 galaxy at redshift 0.00973. Many observers have found a deficit of counts under its power-law X-ray continuum in the soft X-ray range, which they interpret as photoelectric absorption by ionised gas in our line of sight to the active nucleus of the galaxy (e.g. George et al. \\cite{george1998}) - a warm absorber. It has also been claimed, alternatively (Ghosh et al. \\cite{ghosh}), that there is an excess of counts over the power law in the soft X-ray range. De Rosa et al. (\\cite{derosa}), using data from a $\\sim$ five day observation with BeppoSAX in 1998, detect both a soft excess and warm absorption in their soft X-ray spectrum. They also found a high-energy cut-off of the power-law continuum at 340 $^{\\rm +560}_{\\rm -107}$~keV.  With the advent of XMM-Newton and Chandra it has become possible, due to the high resolution of the X-ray spectrometers carried on these missions, to study warm absorbers in unprecedented detail, as was evident after the first high resolution spectrum of a Seyfert galaxy became available (Kaastra et al. \\cite{kaastra2000a}). The first high resolution X-ray spectrum of NGC 3783 (Kaspi et al. \\cite{kaspi2000}) was taken in January 2000 using the Chandra HETGS, with an exposure time of 56 ks, and showed a large number of narrow absorption lines from the H-like and He-like ions of O, Ne, Mg, Si, S and Ar, as well as L-shell transitions of \\ion{Fe}{xvii}$-$\\ion{Fe}{xxi}, and a few weak emission lines mainly from O and Ne. The blueshifts of the absorption lines indicated that the warm absorber was outflowing at -440 $\\pm$ 200 km s$^{\\rm -1}$. Further analysis of this dataset (Kaspi et al. \\cite{kaspi2001}) also confirmed the presence of \\ion{Fe}{xxii} and \\ion{Fe}{xxiii}, and revised the estimate of the absorber's blueshift to -610 $\\pm$ 130 km s$^{\\rm -1}$; the emission lines were at the systemic velocity of the galaxy. Using regions of the continuum where line absorption is not present, a continuum model was fitted, which required deep \\ion{O}{vii} and \\ion{O}{viii} absorption edges (implying an N$_{\\rm H}$ of order 10$^{\\rm 22}$ cm$^{\\rm -2}$). The absorption and emission in the gas were then modelled using photoionisation calculations. The model used involved two phases of gas at different levels of ionisation (with an order of magnitude difference between the ionisation parameters) and with different global covering factors.  NGC 3783 is known to show UV absorption lines intrinsic to the active nucleus (Maran et al. \\cite{maran}), and there is much interest in attempting to connect the X-ray and UV warm absorbers in this object (e.g. Shields \\& Hamann \\cite{shields}). Kaspi et al. (\\cite{kaspi2001}) concluded that the lower-ionisation component in their warm absorber model could give rise to the UV absorption observed in this source, but that this was highly sensitive to the unobservable UV-to-X-ray continuum.  Kaspi et al. (\\cite{kaspi2001}) also obtained the first high resolution spectrum of the region around the Fe K$\\alpha$ emission line, as the HETGS range extends to these energies. They showed that the Fe K$\\alpha$ line was narrow, unresolved even at the HETGS resolution, implying that it originated in the torus region of the AGN. Nandra et al. (\\cite{nandra}) had previously fitted a relativistically broadened Fe K$\\alpha$ line to ASCA spectra of NGC 3783, and the De Rosa et al. (\\cite{derosa}) analysis of BeppoSAX data detects the presence of both broad and narrow components to the line.  The most recent Chandra paper on NGC 3783 (Kaspi et al. \\cite{kaspi2002}) presents a 900 ks exposure obtained from several observations between February and June 2001, which produced an outstandingly high signal-to-noise HETGS spectrum, the best yet published from an AGN source. H-like and He-like Ca, H-like C, various more lowly ionised states of Si and S, iron species probably up to \\ion{Fe}{xxv} and cooler M-shell iron are added to the list of ions detected. The mean blueshift relative to the systemic velocity was found to be -590 $\\pm$ 150 km s$^{\\rm -1}$, with the absorption lines of many ions being resolved to have a mean FWHM of 820 $\\pm$ 280 km s$^{\\rm -1}$. The profiles of the absorption lines show asymmetry which, at least in the case of \\ion{O}{vii}, originates from the presence of two absorbing systems whose velocity shift and FWHM are consistent with those identified in the UV. Evidence has also been found from this dataset by Behar \\& Netzer (\\cite{behar2002}), using the inner-shell absorption lines of silicon, that the warm absorber has a continuous distribution of ionisation parameters, and not just two discrete phases. The narrow Fe K$\\alpha$ line is now resolved to have a FWHM of 1720 $\\pm$ 360 km s$^{\\rm -1}$, and it is accompanied by a Compton shoulder redwards of the line itself, though still by no relativistically (or otherwise) broadened component.  This paper describes the results of work on a 40 ks observation of NGC 3783 with the XMM-Newton Observatory. We obtained UV imaging and photometry from the Optical Monitor (OM; Mason et al. \\cite{mason}) simultaneously with X-ray data from the EPIC-PN (Str\\\"{u}der et al. \\cite{struder}) and RGS (den Herder et al. \\cite{denherder2001}) instruments. The EPIC cameras, although they only have CCD spectral resolution, have a far larger effective area than HETGS around 6.4 keV where Fe K$\\alpha$ emission is observed, requiring far shorter exposures to gain high-statistics spectra at these energies. Again, although the spectral resolution of the RGS is slightly less than that of the HETGS, the combined effective area of the RGS is several times that of the HETGS over much of the energy band where they are both sensitive. Also, the RGS bandpass extends up to about 38~${\\rm \\AA}$ (0.326~keV), whereas the HETGS effective area falls off quickly past 23~${\\rm \\AA}$ (0.539~keV), so we can make important additions to the large amount of knowledge already gained on NGC 3783 by Kaspi et al. (\\cite{kaspi2000}, \\cite{kaspi2001} and \\cite{kaspi2002}).  Partly motivated by the lower statistical quality and resolution of our current RGS spectrum, which makes it harder to disentangle the many absorption features displayed in the 900 ks Chandra dataset, we have concentrated on self-consistent global fitting of the warm absorber rather than detailed measurements of the individual absorption lines. The RGS spectrum of NGC 3783 stands on the threshold between traditional X-ray astronomy - where spectroscopy relied on the $\\chi$$^{\\rm 2}$ fitting of models to often low spectral resolution and poor statistics data - and traditional optical astronomical spectroscopy, where the instrumental resolution and collecting area are sufficiently good to allow precision measurements of individual spectral features. Analysing an RGS spectrum as complex as this requires far more careful usage of $\\chi$$^{\\rm 2}$ fitting techniques than with spectra of lower resolution; there are so many datapoints, and so many interlinked parameters in the increasingly complex spectral models which are fitted to them, that the practical considerations in the global modelling of such spectra will become increasingly important to X-ray astronomers who work on high-resolution AGN spectroscopy. In this paper, within our analysis of the RGS spectrum of NGC 3783, we attempt to define a practical methodology for the modelling and fitting of such datasets.  This work is a prelude to the analysis of a 280 ks XMM-Newton observation of NGC 3783, which has now been carried out, that will provide an RGS spectrum with equivalent signal-to-noise to the 900 ks HETGS spectrum. This new dataset will extend to higher wavelengths the range which has already been observed by the HETGS, and, since it will be accompanied by high time resolution UV photometry by the OM, will be used to investigate further the multiwavelength variability of this AGN. The long observation will also enable time-resolved spectroscopy of the Fe K$\\alpha$ line region. ",
        "conclusions": "\\subsection{Ultraviolet image}  The XMM OM provides us with the first ever UV image of the nuclear regions of NGC 3783. Although the active nucleus is by far the brightest source of ultraviolet light in the galaxy, our image also shows UV emission from the spiral arms, which presumably traces the location of star formation activity. Although the spiral arms are visible in UV, the bar of the galaxy (oriented approximately North-to-South in our image) is not strongly present, indicating that star formation is much less important in the bar than it is in the arms, or alternatively that dust is obscuring the UV emission from the bar.  \\subsection{Lightcurves}  It is clear from Fig.~\\ref{ltcv} that there is unlikely to be sufficient spectral variability in this source to have a serious effect on our modelling of the continuum underlying the RGS spectrum. However, there is clearly variability in the flux and spectral shape on many different timescales. Intensity variations on the scale of about one to three hours, which we see most prominently in the 0.2$-$2~keV band, seem to modulate both soft and hard bands in the same way. Also, the changes that we observe in the hardness ratio demonstrate variability in the relative intensities of the two bands over the course of a few hours. Fig.~\\ref{colour} indicates that the spectrum generally softens as the source intensity increases, although at the highest intensity point, that is at the very end of our observation, the source is actually in a harder state. The UV lightcurve is almost completely flat, within the errors, implying less than 10\\% variability in the UV flux from the nucleus during our observation. There is therefore no evidence in our dataset for a direct correlation between UV and X-ray variability over the course of a few hours.  \\subsection{Continuum and iron K$\\alpha$ line}  We find that the 0.2$-$10~keV spectral continuum of NGC 3783 is dominated by a power-law at high energies, whilst there is significant absorption by ionised gas - the warm absorber - at lower energies. There is no evidence for a soft excess, although of course the presence of a soft excess could be masked by the warm absorber. There is also no requirement for a reflection component in our spectrum.  The iron K${\\alpha}$ emission line has two components; one narrow component at 6.40 keV, and a broader component which is redshifted by about 5000 km s$^{\\rm -1}$. The profile of this composite line looks superficially similar to that of the NGC 3783 spectrum (1) of Nandra et al. (\\cite{nandra}), although our broad component is far narrower than the one they observed. We were able to fit our broad component with a gaussian; it does not contain enough flux to obtain good constraints on a relativistic disc line model. Our narrow Fe K${\\alpha}$ line is apparently similar to that observed by Kaspi et al. (\\cite{kaspi2002}), although our derived parameters are not well constrained.  The broad component of the Fe K${\\alpha}$ line implies emission from gas that is infalling close to the nucleus. The line broadening would presumably be dynamic rather than thermal, as Fe K${\\alpha}$ arises from fluorescence of neutral iron, so the emitting matter is in a region of high speed motion under strong gravity. Certainly, the derived velocity broadening of this component, at 11000 $\\pm$ 4000 km s$^{\\rm -1}$, would place it in the very broad line region of Winge et al. (\\cite{winge}). We were unable to ascertain whether the broadening is actually relativistic, although if we were dealing with emission from close enough to the black hole to be relativistically broadened, we might expect the iron to be rather more highly ionised than the energy of this line implies.  That the broad component of the Fe K$\\alpha$ line was not observed by Kaspi et al. (\\cite{kaspi2002}) in the 900 ks Chandra observation, where the signal-to-noise in this spectral region was very good, may imply that the spectrum has changed between the two observations. Infalling gas close to the central engine of the AGN, but not yet close enough to become too hot to emit Fe K${\\alpha}$, might be a transitory phenomenon which is not always observable.  The Fe K${\\alpha}$ emission is accompanied by another line at higher energy, which is probably a combination of Fe K${\\beta}$ and \\ion{Fe}{xxvi}. The apparent broadening and redshift of this line are, then, a function of the flux and redshift of the components that it contains. \\ion{Fe}{xxvi} emission would not originate from the same material as Fe K${\\alpha}$; instead, this must be emission from much hotter gas.  \\subsection{Warm absorber}  \\subsubsection{Modelling strategy}  We developed a strategy for modelling the complex RGS spectrum of this object using SPEX, with the aim of reproducing both line and continuum absorption self-consistently, and of obtaining some physical insight into the properties of the warm absorber.  This process has several steps. Firstly, \\emph{slab} - a model which applies absorption by individual ions in a photoionised medium to a given spectral continuum - is used to identify the patterns of line absorption by different ionic species in the warm absorber, and to measure the blueshift of the lines. This model is also used to investigate the contributions that each ion can make to continuum absorption. Once the lines have been identified, the relative strengths of line absorption by the different iron species are used to derive the ionisation structure of the absorber. Iron is used for this purpose due to its wide range of ionisation states, several of which can be observable in gas at a given level of ionisation, thus providing a good constraint on the ionisation parameter of the medium. We contend that this gives a more reliable picture of the ionisation structure of the absorber than a simple measurement of \\ion{O}{vii} and \\ion{O}{viii} edge opacities, which is dependent on continuum absorption (a more ambiguous diagnostic than line absorption) from only two ions.  At this stage, the \\emph{xabs} model, which applies line and continuum absorption by a given column of photionised gas at a given ionisation parameter to the spectral continuum, is used to fit an overall column for the absorber at the ionisation parameter independently derived from the iron absorption. Multiple \\emph{xabs} models can be used to represent the different ionisation phases within the absorber. This is done with the assumption of solar abundances within the absorber, but, once an estimate of the overall column has been derived, the individual elemental abundances can be allowed to vary. To get columns for the individual ionic states, one can return to \\emph{slab} to fit values for these.  At every stage of the process outlined here, due to the likelihood of \\lq blind\\rq \\, $\\chi$$^{\\rm 2}$ fitting being misled by the enormous complexity of the spectrum and model, it is necessary to spend some time getting to know the model properties by trial and error. A probable range of realistic values needs to be established before using $\\chi$$^{\\rm 2}$ minimisation to home in on the final value for a parameter, being careful to ignore any spectral regions that will bias the fit towards unphysical results.   \\subsubsection{Two-phase warm absorber model}  The two-phase warm absorber model reproduces the global form of the spectrum very well. It indicates that the ratios of elemental abundances to each other in the high-ionisation phase of the gas are probably not too different from solar. However, in order to fit the very deep UTA associated with the low-ionisation phase of the gas, iron would be ten times more abundant, relative to its solar value, than oxygen. This is sensitive to the turbulent velocity and overall absorbing column assumed for the low-ionisation phase, and the level of the spectral continuum itself. If the spectral continuum was higher (especially at the long wavelength end of the spectrum, where continuum absorption by \\ion{C}{iv} and \\ion{C}{v} becomes important), the equivalent hydrogen column of the low-ionisation phase would need to be higher and thus the iron abundance could be a lot lower. The value of $\\xi$ in the model is too high for some of the lower charge states of Fe, therefore requiring high Fe abundance to compensate for the low fractional ionic abundances that the model produces. The low-ionisation phase might also have a wider distribution of $\\xi$. Moreover, the ionisation balance calculations of Fe-M are very uncertain; in particular, the dielectronic recombination rates for these ions have never been measured or calculated with modern codes.  The much shallower UTA in IRAS 13349+2438 (Sako et al. \\cite{sako}) requires an iron abundance 2$-$3 times that of oxygen, and the UTA in the LETGS spectrum of NGC 5548 (Kaastra et al. \\cite{kaastra2002b}) also suggests an overabundance of iron. It is possible that these apparent high iron abundances could be due to a covering factor effect (e.g. Arav et al. \\cite{arav}). The iron abundance in the UTA is effectively measured against the abundances of C, N and O, and if these lines are highly saturated their equivalent widths will be determined by the covering factor and velocity width, and will be only weakly dependent on the column density.  The position of the UTA also coincides with the \\ion{O}{vii} absorption edge (16.78~$\\rm \\AA$ (0.7389~keV) in the rest frame), and the deeper this edge, the lower the iron abundance required to fit the UTA. We would not expect significant \\ion{O}{vii} absorption to originate from the low-ionisation phase itself, as it is not highly ionised enough. The high-ionisation phase, on the other hand, could be a plausible source of \\ion{O}{vii} edge absorption. At the derived ionisation parameter of this phase, though, oxygen will be predominantly in the form of \\ion{O}{viii}. Certainly, the fitted equivalent hydrogen column of the high-ionisation phase, at 2.8 x 10$^{\\rm 22}$ cm$^{\\rm -2}$, is easily high enough to give rise to significant edge absorption. Examination of the detailed form of the spectrum where the edges are expected, however, shows that any such edges would be completely masked by L-shell and M-shell iron absorption. We propose, then, that although \\ion{O}{vii} and \\ion{O}{viii} edge opacity is certainly part of the picture, absorption by L-shell and M-shell iron is likely to be more important in determining the detailed form of the spectrum in these regions.  The \\ion{O}{viii} absorption is, in fact, seriously overpredicted in the two-phase model (N$_{\\rm \\ion{O}{viii}}$ $\\geq$ 10$^{\\rm 18}$ cm$^{\\rm -2}$), implying a very deep \\ion{O}{viii} Ly$\\alpha$ absorption line (at 18.97~$\\rm \\AA$ (0.6536~keV), see Fig.~\\ref{rgs_spec}) that is not matched by the data. It is of course possible that this line is partially filled in by \\ion{O}{viii} line emission, but we cannot test this with the current dataset. It can also be seen from Fig.~\\ref{rgs_spec} that the iron absorption is overpredicted at the low-wavelength side of the 15~$\\rm \\AA$ feature and increasingly underpredicted at lower wavelengths. Fig.~\\ref{model_comps} shows that the two-phase model compensates for this by generating neon edge absorption at $\\sim$ 9.5~$\\rm \\AA$. This clearly demonstrates that it was not correct to approximate log $\\xi$ of the high-ionisation phase with a single value of 2.4 - there is a lot of gas which is more highly ionised than this, so the high-ionisation phase should be more accurately modelled as containing a range of ionisation parameters from 2.4 upwards. The low-ionisation phase, too, is not perfectly explained by a single phase with log $\\xi$ of 0.3, but the single ionisation parameter is a good approximation in this case.  A comparison between the parameters of our two-phase model, and that of Kaspi et al. (\\cite{kaspi2001}), is given in Table~\\ref{modelcomp}. Their low-ionisation phase corresponds most closely to our high-ionisation phase, whilst their high-ionisation phase is much more highly ionised than ours. It probably corresponds to the highly-ionised iron at the low-wavelength end of our spectrum which our two-phase model does not explain. Kaspi et al. (\\cite{kaspi2001}) do not model our low-ionisation (UTA) phase, although they report the presence of the UTA. The warm absorber modelled by De Rosa et al. (\\cite{derosa}) is consistent with the low-ionisation component of Kaspi et al. (\\cite{kaspi2001}), and thus with our high-ionisation component.  The two models, although different, are not in fundamental disagreement with each other. The differences between them may be due to the methods used to model the spectrum. The Kaspi et al. (\\cite{kaspi2001}) modelling relies on the use of measured equivalent widths of the absorption lines as well as the fitting of a continuum model to \\lq line-free zones\\rq \\, where absorption lines are not expected. We, on the other hand, use the line absorption as the prime diagnostic and tie the continuum absorption directly to this. The elemental abundances in their model were kept at solar values, whilst ours were allowed to vary. These two phase models, though, are just approximations to a multi-phase absorber, as we discuss above, and as already pointed out in Behar \\& Netzer (\\cite{behar2002}).  \\begin{table}  \\caption[]{Comparison of the properties of the two-phase warm absorber model presented here with that of Kaspi et al. (\\cite{kaspi2001})} \\label{modelcomp} $$ \\begin{array}{p{0.9in}p{0.5in}p{0.9in}p{0.7in}} \\hline \\noalign{\\smallskip} Phase & Property & Our model & Kaspi et al. \\cite{kaspi2001} model \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} High-ionisation & log $\\xi$$^{\\mathrm{a}}$  & 2.4 & 3.5 \\\\ High-ionisation & N$_{\\rm H}$$^{\\mathrm{b}}$  & 2.8 $^{\\rm +0.01}_{\\rm -0.3}$ x 10$^{\\rm 22}$ & 1.6 x 10$^{\\rm 22}$ \\\\ Low-ionisation & log $\\xi$$^{\\mathrm{a}}$  & 0.3 & 2.5 \\\\ Low-ionisation & N$_{\\rm H}$$^{\\mathrm{b}}$  & 5.4 $^{\\rm +0.05}_{\\rm -0.5}$ x 10$^{\\rm 20}$ & 1.6 x 10$^{\\rm 22}$ \\\\ \\noalign{\\smallskip} \\hline \\end{array} $$ \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Log of the ionisation parameter, where $\\xi$ is in erg cm s$^{\\rm -1}$ \\item[$^{\\mathrm{b}}$] Equivalent hydrogen column of phase in cm$^{\\rm -2}$ \\end{list}{}{} \\end{table}  \\subsubsection{Detailed modelling}  Our detailed model, using \\emph{slab}, provides a list of the total columns of the main ions present in the multi-phase warm absorber, summed across all phases. This model (Fig.~\\ref{spec_1} and Fig.~\\ref{spec_2}) provides a better reproduction of the detailed form of the spectrum than the simple \\emph{xabs} two-phase approximation. The agreement of our model with the data in the region of L-shell iron absorption, in particular, is excellent in some ranges. There are still some significant residuals in this region, though. The wavelengths and strengths of the L-shell iron absorption lines and blends are very well known, so it is unlikely that this is the source of the discrepancy. Some of the absorption may originate from elements, particularly calcium, which are not yet included in our model. It is possible to reduce some of the residuals by increasing the \\ion{O}{viii} continuum absorption, although this then introduces very deep \\ion{O}{viii} absorption lines which do not appear in our data. One interesting explanation for these residuals might be that our model lacks the higher level transitions associated with the UTA (Behar et al. \\cite{behar}) which appear when there is a high column density in the UTA ions.  The derived columns of the L-shell iron states above \\ion{Fe}{xx} do indicate that there is a lot of gas in the hot phase of this warm absorber at log ionisation parameters above 2.4, as shown in Fig.~\\ref{col_xi}. This may be part of the explanation for the inability of the two-zone model to accurately reproduce the \\ion{O}{viii} absorption; much of the gas in the hot phase may simply be too highly ionised for \\ion{O}{viii} to form in significant quantities. The derived columns of \\ion{O}{vii} and \\ion{O}{viii} ($\\sim$ 10$^{\\rm 17}$ cm$^{\\rm -2}$) in this detailed model are too low to produce significant edge absorption, but their values are fairly uncertain.  The effective area of the HETGS system falls off above around 23~${\\rm \\AA}$ (0.539~keV), so since the RGS spectrum extends up to ~38 ${\\rm \\AA}$ (0.326~keV), we can add quantitively to the knowledge gained from the Chandra spectrum of NGC 3783. In this high wavelength range we see absorption lines from \\ion{N}{vi}, \\ion{N}{vii}, \\ion{C}{vi} and possibly \\ion{C}{v}, \\ion{Si}{xi}, \\ion{Si}{xii}, \\ion{S}{xii}, \\ion{S}{xiii}, \\ion{S}{xiv}, \\ion{Ar}{xiii} and \\ion{Ar}{xiv}. Direct superimposition of the model on the spectrum (Fig.~\\ref{spec_1} and Fig.~\\ref{spec_2}) reveals clearly the inaccuracies in the model wavelengths of various lines, and shows that there are further absorption lines in the spectrum that are yet to be included in our model. This work will be carried forward using the 280 ks RGS exposure of NGC 3783.  What is the connection between the higher and lower ionisation phases of the warm absorber? They appear to be flowing towards us at a similar speed, around 800 km s$^{\\rm -1}$. The high-ionisation phase has an equivalent hydrogen column about fifty times that of the low-ionisation phase in our line of sight, and the low-ionisation phase apparently has a very high iron abundance. The spread of values of the ionisation parameter in the two phases could give us a clue as to what is going on. As indicated by the results of the \\emph{xabs} modelling, and also in Fig.~\\ref{col_xi}, the low-ionisation phase appears to be concentrated around a single ionisation parameter, whilst the high-ionisation phase contains a range of levels of ionisation. When the central engine of NGC 3783 is radiating at a given luminosity, the ionisation parameter of the absorbing outflow is determined by the density and distance from the source of the absorbing gas. $\\xi$ is more sensitive to changes in distance than changes in density. One could, then, surmise that the wide range of $\\xi$ in the high-ionisation phase implies absorption by gas at a wide range of distances and (to a lesser extent perhaps) densities, whereas the low-ionisation, UTA phase absorption is quite localised to gas at a certain distance from the source, and at a particular density. This would imply the existence of a highly-ionised, multi-temperature outflow (as in the multi-temperature wind model posited by Krolik \\& Kriss \\cite{krolik}), spread over a range of distances from the source, and denser and perhaps iron-rich material at a fixed location within the circumnuclear environment.  Sako et al. (\\cite{sako}) claim that the low-ionisation phase of the warm absorber in IRAS 13349+2438 might originate in the dusty torus surrounding the active nucleus, which would fit in with the scenario described here. Perhaps the abnormally high iron content of the corresponding phase in NGC 3783 is due to the sublimation of iron-rich dust from the edge of a dusty torus. If this was the case, the high-ionisation phase would form the bulk of the warm absorber, and wherever it encountered the edge of the dusty torus it would sweep up denser material; this would then be photoionised and the dust it contained would be sublimated, giving rise to the high iron content of the low-ionisation phase. This places constraints, of course, on the composition of the dust, which would have to be justified. The higher signal-to-noise RGS spectrum which has now been obtained will allow us to investigate the relative abundance of other non-volatile elements which can be bound into dust grains (including silicon and carbon), to test this hypothesis further."
    },
    "0206/astro-ph0206066_arXiv.txt": {
        "abstract": "{We present results of spectroscopic observations of selected stars in the southern open cluster NGC~6134. We have determined the rotational velocities of the six known \\dss\\ in NGC~6134 as well as several other non-variable stars with similar colour temperature in order to investigate if \\vsini\\ and variability is somehow connected: we find no such correlation. We also compare the distribution of \\vsini\\ of \\dss\\ and  non-variable stars with four other well-studied open clusters to look for any systematic behaviour, but we find no conclusive evidence for \\vsini\\ and variability to be connected. We have also used the spectra to carry out an abundance analysis of the \\dss\\ in NGC~6134 to confirm the high metal content of the cluster. We find $[{\\rm Fe/H}] = +0.38\\pm0.05$ which is in agreement with the result obtained from Str\\\"omgren photometry. We also present $\\Delta$a photometry of the cluster, but we find no chemical peculiar stars based on this index. ",
        "introduction": "}  The open cluster NGC 6134 has six confirmed \\ds\\ members (Frandsen et al. 1996) of which five are multiperiodic. High quality Str\\\"omgren photometry has been carried out by Bruntt et al.\\ (1999), except that the $c_1$ index (requiring $u$ images) could not be obtained with the CCD detector. This paper describes a continuation of the earlier programs. We use spectroscopy and $\\Delta$a-photometry (Maitzen 1976, Maitzen \\& Vogt 1983) to obtain more parameters and additional information about the stars in the cluster. These parameters are needed for the modeling and interpretation of the oscillations in the $\\delta$ Scuti stars. Here we will derive rotational velocities (\\vsini) for variable and non-variable stars and carry out abundance analysis for the \\dss.   \\begin{figure*}\\centering \\begin{center} \\subfigure[The non-stellar background signal is very high in this case\\label{subfig:blowup}] {\\includegraphics[width=\\dplotwidth, bb=110 256 476 508, clip]{MS2364f1a.eps}} \\qquad \\subfigure[For this star the background level is still significant but not dominating]{\\includegraphics[width=\\dplotwidth, bb=110 256 476 508, clip]{MS2364f1b.eps}} \\caption{Expanded plot of the cross section of an order for two different stars. The $x$-axis measures the distance from the center of the order in pixels. The horizontal bars in the lower part of each plot show how we have set the background regions in IRAF during manual editing of the apertures when using the task \\task{apall}. In \\ref{subfig:blowup} the high background signal indicates that the Moon is nearby\\label{backgr}} \\end{center} \\end{figure*}   In principle, with the combination of the information from the frequencies of the \\dss\\ and the constraints imposed on the variables by being members of an open cluster we will be able to make a very powerful test of stellar evolution. The analysis of \\dss\\ is very complicated in stars that are fast rotators, so it is important to locate slow rotators before extensive campaigns are started to obtain oscillation spectra with very low noise levels, leading to the detection of many oscillation frequencies. For example, the open cluster Praesepe has 14 \\dss\\ and all but one rotate fast ($v \\sin i > 100$~\\kms) (Rodr\\'{\\i}guez et al.\\ 2000).         % This has complicated the determination of the stellar parameters considerably and made the identification of the observed frequencies with radial or non-radial modes very difficult (P\\'erez Hern\\'andez et al.\\ 1999, Kjeldsen et al.\\ 1998). The situation is not quite as bad for the Pleiades and Hyades clusters where a few slowly rotating \\dss\\ are found.  The nature of the excitation mechanism in \\dss\\ is an unsolved problem -- especially the understanding of which among the many possible modes predicted from models have obserable amplitudes. We have looked for systematic differences between variable and non-variable stars. When the variables were discovered, a large set of non-variable stars (amplitude $< 1$~mmag) were identified. It is therefore possible to find pairs of stars, variable and non-variable (twins), with similar photometric parameters. The \\dss\\ are located in the instability strip and intrinsically unstable due to the opacity mechanism; but the \\dss\\ have modes that do not seem to follow any regular pattern. Many stars in the instability strip do not vary at all. Perhaps the explanation is the difference in observables like \\feh\\ and \\vsini\\ among \\dss\\ and non-variable stars in the instability strip. We will examine this hypothesis by comparing variable and non-variable stars with approximately the same colour-temperature and evolutionary state.    \\begin{figure*}\\centering \\includegraphics[width=\\textwidth]{MS2364f2.eps} \\caption[plot_ha.pro]{\\halpha\\ of the six \\dss\\ plotted for comparison. The dashed line at 0.4 is just a visual aid\\label{spectra}} \\end{figure*}    Finally, we try to confirm the metallicity determined from Str\\\"omgren photometry (Bruntt et al.\\ 1999) and to locate possible Am or Ap-type stars among the $\\delta$ Scuti stars. We will do abundance analysis of the \\dss\\ and we also investigate any possible chemical peculiarity with $\\Delta$a photometry. It seems that the diffusion leading to chemically peculiar stars is coupled with slow rotation and seems to exclude variability (Kurtz 2000). But since NGC~6134 is a metal rich cluster, with a large number of A-type stars with high values of the Str\\\"omgren $m_1$ index -- also in case of the \\dss\\ -- it would be interesting if some of these are CP-type stars.  The observations are described in Section~\\ref{sec_obs} and in Section~\\ref{sec_red} we describe the reduction of the spectra. The calibration of the stellar parameters is given in Section~\\ref{sec_cal} and in Section~\\ref{sec_rot} we determine the rotational velocities. The abundance analysis is carried out in Section~\\ref{sec_abu}. The data reduction and results of the $\\Delta$a photometry is given in Section~\\ref{sec_del}. In Section~\\ref{sec_one} we gather in information on each of the known \\dss\\ and in Section~\\ref{sec_dis} we give a general discussion of our results before giving our conclusions in Section~\\ref{sec_con}. ",
        "conclusions": "}   We have presented the results of our analysis of spectra of variable and non-variable stars in the metal rich open cluster NGC~6134. We have determined the rotational velocity of several stars and made abundance analysis of the \\dss\\ stars. We summarize our conclusions below:  \\begin{itemize}  \\item{We have determined \\vsini\\ for the six known \\dss\\ and a set of known non-variable A-type stars in NGC 6134}  \\item{We have determined \\feh\\ from spectroscopy and find a good agreement with the value from the $m_1$ index. We confirm that NGC~6134 is quite metal rich with \\feh\\ around $0.3\\pm0.1$ (error estimate includes systematic errors). From the $m_1$ index some stars seem to have enhanced atmospheric metal content and could be Am-type stars. We find that one of the \\dss\\ seem to be a marginal Am-type star (star \\#508)}  \\item{We have presented $\\Delta$a photometry which was carried out in order to detect any chemically peculiar stars. We find no classical Am-type stars in our sample of stars but we note that we did not obtain $\\Delta$a for the interesting \\ds\\ \\#508}  \\item{The distributions of rotational velocities of \\dss\\ and non-variable A-stars in five well studied open clusters have been analysed: we find that the distributions are quite different but no safe conclusions can be drawn due to the small number of stars. For NGC~6134 we find that half of the known \\dss\\ have low \\vsini\\ and they would be suitable for modeling by the use of asteroseismology in the future}  \\end{itemize}"
    },
    "0206/astro-ph0206299_arXiv.txt": {
        "abstract": "We investigate star formation along the Hubble sequence using the ISO Atlas of Spiral Galaxies.  Using mid-infrared and far-infrared flux densities normalized by K-band flux densities as indicators of recent star formation, we find several trends.  First, star formation activity is stronger in late-type (Sc~- Scd) spirals than in early-type (Sa~- Sab) spirals.  This trend is seen both in nuclear and disk activity. These results confirm several previous optical studies of star formation along the Hubble sequence but conflict with the conclusions of most of the previous studies using IRAS data, and we discuss why this might be so.   Second, star formation is significantly more extended in later-type spirals than in early-type spirals.  We suggest that these trends in star formation are a result of differences in the gas content and its distribution along the Hubble sequence, and it is these differences that promote star formation in late-type spiral galaxies.  We also search for trends in nuclear star formation related to the presence of a bar or nuclear activity.  The nuclear star formation activity is not significantly different between barred and unbarred galaxies.  We do find that star formation activity appears to be inhibited in LINERs and transition objects compared to H~{\\small II} galaxies.  The mean star formation rate in the sample is 1.4~M$_\\sun$~yr$^{-1}$ based on global far-infrared fluxes.  Combining these data with CO data gives a mean gas consumption time of 6.4~$\\times$~10$^{8}$~yr, which is $\\sim$ 5 times lower than the values found in other studies.  Finally, we find excellent support for the Schmidt Law in the correlation between molecular gas masses and recent star formation in this sample of spiral galaxies. ",
        "introduction": "\\subsection{Scientific Background}  Spiral galaxies have been recognized as sites of on-going star formation since the definition of the Hubble sequence \\citep{h26}.  However the relation between star formation activity and the location on the Hubble sequence has been a continuing controversy for the past twenty years.  Early optical studies comparing star formation activity in spiral galaxies included \\citet{s82}, \\citet{kk83}, \\citet{k83}.  They found a general trend in star formation activity with Hubble type, with stronger star formation in Sc spirals than in Sa spirals. However, \\citet{setal83}, using ground-based 10~$\\mu$m observations, demonstrated that star formation activity was not correlated with any global galaxy property, including Hubble type.  This laid the groundwork for a debate on trends in star formation along the Hubble sequence that continues to the present.  With the advent of IRAS, several new studies on this theme were published, including \\citet{detal84}, \\citet{dbs87}, \\citet{pr89}, \\citet{dy91}, \\citet{if92}, \\citet{tts96}, and \\citet{dh97}.  Most IRAS studies found no trend in star formation along the Hubble sequence, with the exceptions being \\citet{dbs87} and \\citet{pr89}, who found that Sa galaxies are generally weaker sites of star formation than Sc galaxies, and \\citet{dh97}, who found a population of Sa galaxies with relatively enhanced star formation activity.  But optical and ultraviolet studies on the problem also continued, with \\citet{detal87}, \\citet{detal94}, and \\citet{ktc94} all finding trends in increasing star formation along the Hubble sequence from Sa to Sc galaxies.  To confuse matters further, \\citet{yetal96} found trends in H$\\alpha$ surface brightness along the Hubble sequence but they argued that the trend did not necessarily reflect a true trend in star formation activity.  Infrared spectroscopy of the C~{\\small II} 158~$\\mu$m line performed with ISO by \\citet{lvhetal99} also found a trend in increasing star formation from early- to late-type spiral galaxies, but this survey had a small sample of only 19 spiral galaxies.  Clearly, opinion is sharply divided on star formation activity along the Hubble sequence.  Some of the reasons for these discrepant results may include the following.  1) Extinction is always a concern in optical or ultraviolet studies, since star formation regions are always dusty.  2) The IRAS sweep-scan technique is not ideal for measuring the flux from faint, extended sources.  For faint, extended galaxies the IRAS point source algorithm may mistake the faint outer regions of the galaxy itself for the background and therefore undermeasure the flux.  3) To normalize for galaxy mass, most IRAS studies divide the far-infrared fluxes by B-band fluxes. Since the B-band fluxes represent the fluxes from the youngest, bluest stars in any galaxy, this tends to cancel out any trend in star formation activity in the far-infrared data.  4) The morphological classifications used are sometimes biased towards star formation.  In particular, the Revised Shapley Ames (RSA) Catalog \\citep{st87} uses star formation activity in its classification scheme \\citep{sb94}, so results on star formation trends found by any paper using the RSA classifications may merely reflect the classification criteria.  \\subsection{The Design of this Study}  As a fresh approach to this question, we have completed a new survey of spiral galaxies with the Infrared Space Observatory (ISO) \\citep{ketal96} at mid-infrared and far-infrared wavelengths, supplemented with  ground-based JHK photometry. The data are presented in the ISO Atlas of Spiral Galaxies \\citet{betal02}  (henceforth referred to as Paper 1).  This study incorporates several features that address the concerns mentioned above.  1) We minimize uncertainties related to extinction by working strictly with infrared data.  2) The ISO data has several advantages over IRAS data.  ISO has improved sensitivity and angular resolution that enables us to  distinguish nuclear from disk emission. Furthermore, ISO operates in a point and integrate mode, and the background is measured separately from locations pre-selected from IRAS cirrus maps, so the flux from faint extended galaxies can be measured more accurately (although the flux measurements are still limited by background structure). 3) We normalize mid-infrared and far-infrared fluxes with K-band fluxes, which represent more accurately the total stellar population than B-band fluxes.  4) We use the Third Reference Catalogue of Bright Galaxies (RC3; \\citet{detal91}) for galaxy classification, since the RC3 classifications are strictly dependent on morphological features such as the bulge to disk ratio and the tightness of the spiral arms \\citep{d59}.  We first summarize the sample selection, observations, and data reduction, the details of which can be found in Paper 1. Next, we discuss our diagnostics for studying star formation activity.  Then we present our results on trends in star formation along the Hubble sequence, and we compare these results to previous results and discuss the underlying mechanisms for producing the trends.  Next, we examine the galaxies for any possible trends in nuclear star formation related to the presence of a bar or nuclear activity.  We conclude by calculating quantitative star formation rates and gas consunsumption times and by comparing these results to those in the literature. ",
        "conclusions": "Using the ISO Atlas data for an optically-selected, magnitude-limited sample of spiral galaxies, combined with normalization by K-band flux densities, we have examined the evidence for star formation in three regions of these galaxies: integrated over the entire 135~arcsec of the galaxies (global), the central 15~arcsec (nuclear), and in the disk itself excluding the nuclear region. We have found strong trends in star formation activity along the Hubble sequence in all three regions with statistical significance of $\\sim$~3-6~$\\sigma$. We also show that the distribution of star formation regions is more spatially-extended in late-type galaxies than early-type galaxies.  These results confirm most previous optical and ultraviolet investigations of star formation along the Hubble sequence, but do not confirm the findings of an equally large number of such studies using infrared (mostly IRAS) data, which have found no correlations with Hubble type.  We suggest that the superior angular  resolution in the ISO Atlas data, combined with use of K-band flux densities for normalization, has resolved this long-standing discrepancy between optical and infrared studies.  We argue that the differences in star formation activity along the Hubble sequence may be understood primarily in terms of the differences in the molecular gas density and distribution.  We confirm that star formation is well-correlated with molecular gas mass. We suggest that the strength of the star formation activity increases as the gas density increases along the Hubble sequence. Similarly, we suggest that the spatial extent of star formation activity increases as the spatial extent of the molecular gas increases.  We found no enhancement of nuclear star formation in barred galaxies compared to unbarred galaxies, in contrast with many previous results. This may be because the effect is seen chiefly in early-type galaxies and we have few early-type SA and SAB galaxies in the  sample.  We do find that star  formation seems to be inhibited in LINERs and transition objects compared to H~{\\small II}  galaxies, suggesting that LINER emission is not powered by recent starburst activity  (at least in nearby galaxies).  From the far-infrared fluxes we estimate star formation rates of 1.4~M$_\\sun$~yr$^{-1}$, which is in agreement with previous studies of similar spiral galaxy samples.  However, we find gas consumption times of 6~$\\times$~10$^{-8}$~yr, which is about a factor of 5 lower than most previous studies.  We do find excellent correlation between molecular gas densities and recent star formation activity in accord with the Schmidt Law."
    },
    "0206/astro-ph0206250_arXiv.txt": {
        "abstract": "We discuss {\\em Far Ultraviolet Spectroscopic Explorer} (\\fuse) observations of two early-type stars, DI\\,1388 ($l = 291.3 \\degr$, $b = -41.1\\degr$) and DGIK\\,975 ($l = 287.3\\degr$, $b = -36.0 \\degr$), in the low density and low metallicity ($Z\\sim0.08 Z_\\odot$) gas of Magellanic Bridge (MB). The data have a spectral resolution of about 15,000 and  signal-to-noise ratios range between 10 and 30 per resolution element in the spectra of DI\\,1388 and between 7 and 11 in the spectra of DGIK\\,975. DI\\,1388 is situated near the SMC, while  DGIK\\,975 is closer to the LMC, allowing us to probe the MB gas in a widely different locations. Toward DI\\,1388, the \\fuse\\/ observations show molecular hydrogen,  \\ion{O}{6}, and numerous other atomic or ionic transitions in absorption, implying the presence of multiple gas phases in a complex arrangement. The relative abundance (with respect to \\ion{S}{2}) pattern in the MB along the DI\\,1388 sight line is attributed to varying degrees of depletion onto dust similar to that of halo clouds. The N/O ratio is near solar, much higher than N/O in damped Ly$\\alpha$ systems, implying subsequent stellar processing to explain the origin of nitrogen in the MB. The diffuse molecular cloud in this direction has a low column density and low molecular fraction ($\\log N ({\\rm H}_2) \\approx 15.43$ dex; $f_{\\rm H_2} \\sim 10^{-5} - 10^{-4} $), yet two excitation temperatures ($T_{01} = 94 \\pm\\,^{53}_{27}$ K and $T_{23} = 341 \\pm\\,^{172}_{81}$ K) are needed to fit the distribution of the different rotational levels. Though this is not typically seen in the Galaxy, we show that this is not uncommon in the Magellanic Clouds. H$_2$ is observed in both the Magellanic Stream and the MB, yet massive stars form only in the MB, implying significantly different physical processes between them. In the MB some of the H$_2$ could have been pulled out from the SMC via tidal interaction, but some also could have formed {\\em in situ} in dense clouds where star formation might have taken place. Toward DGIK\\,975, the presence of neutral, weakly and highly ionized species suggest that this sight line has also several complex gas phases. The highly ionized species of \\ion{O}{6}, \\ion{C}{4}, and \\ion{Si}{4} toward both stars have very broad features, indicating  that multiple components of hot gas at different velocities are present. \\ion{C}{4}/\\ion{O}{6} varies within the MB but \\ion{C}{4}/\\ion{Si}{4} is relatively constant for both sight lines.  Several sources (a combination of turbulent mixing layer, conductive heating, and cooling flows) may be contributing to the production of the highly ionized gas in the MB. Finally, this study has confirmed previous results that the high-velocity cloud HVC $291.5-41.2+80$ is mainly ionized composed of weakly and highly ions. The high ion ratios are consistent with a radiatively cooling gas in a fountain flow model. ",
        "introduction": "The Magellanic system is composed of two small irregular galaxies, the Large (LMC) and Small (SMC) Magellanic Clouds, in orbit around the Galaxy. Tidal interactions between these galaxies have produced several high velocity  gas complexes connected to the Clouds \\citep[for a recent study, see][]{putman00}; namely the Magellanic Bridge (MB), the Magellanic Stream, and the leading Arm. In particular for this study, the Magellanic Bridge is a $10\\degr$ region of tenuous gas linking the body of the SMC to an extended arm of the LMC. The formation mechanism responsible for this feature remains unclear, but it is generally agreed that the MB was formed via a tidal encounter between the SMC and LMC. \\citet{gardiner} have produced models that can reproduce simultaneously both the MB and the Magellanic Stream, and find that the MB was most likely pulled from the wing of the SMC 200 Myr ago during a close encounter between the two Clouds. However, its low metallicity \\citep[$Z \\approx 0.08 Z_\\odot$ based on C, N, O, Mg, and Si][]{rol99} does not reflect the current SMC metallicity \\citep[$Z \\approx 0.25 Z_\\odot$, e.g.,][]{russell}, suggesting that the MB gas could be formed from a mixture of SMC gas and an unenriched component or could be much older than the age predicted from theoretical models.  The MB contains early-type stars \\citep{hambly,demers98,rol99}. During their main-sequence lifetimes (as short as 20 Myr), these stars could not possibly migrate from the SMC since they do not exhibit peculiar velocities sufficient to explain their motion over the large distances they would need to cover. Therefore the MB provides the most metal poor gas in our neighborhood to investigate not only the gas, but also to observationally constrain star-formation in low metal and tenuous gas environment. Furthermore, it provides an unique opportunity to study in detail the gas resulting from tidal interactions.  In the Galaxy, star formation generally occurs in dense molecular clouds \\citep[e.g.,][]{evans99}. No direct evidence of molecular clouds have been discovered yet \\citep{smoker00}, although \\citet{kob99} found cold atomic clouds, suggesting molecular condensations.  \\citet{lehner01a} recently reported the results of a program to investigate the chemical composition and abundance pattern of the MB gas toward an early-type star, DI\\,1388, with the {\\em Hubble Space Telescope} {\\em (HST)}\\/ and the Space Telescope Imaging Spectrograph (STIS). The combination of high spectral resolution and high sensitivity in the ultraviolet bandpass ($1150-1730$ \\AA) made it feasible to investigate the chemical composition and the physical conditions within the MB and revealed complex gas phases with neutral gas, weakly and highly ionized gas along the sight line studied. Yet only at lower wavelengths than the STIS bandpass, in the 905--1187 \\AA\\ wavelength range, are the strong hydrogen molecular lines accessible. Other strong atomic and ionic resonance lines are as well only limited in this bandpass, such as the \\ion{O}{6} doublet, a powerful diagnostic of collisionally ionized gas. \\ion{O}{6} is unlikely to be produced by photoionization alone from starlight given that photons with $h\\nu \\ge 114$ eV are needed to convert \\ion{O}{5} to \\ion{O}{6}. Such a bandpass is accessible now with the {\\em Far Ultraviolet Spectroscopic Explorer} ({\\fuse}) and therefore follow-up observations of the low metallicity, tenuous gas of the MB were obtained with this observatory.  In this article, we report on {\\fuse}\\/ observations of two early-type stars in the MB,  DI\\,1388 which is situated approximately mid-way between the SMC and LMC  and DGIK\\,975, which lies at the western end of the Bridge near to the LMC halo. The {\\fuse}\\/ spectra allow us for the first time to make a sensitive and direct search for molecules in the MB. They provide a means to study the collisional gas seen in \\ion{O}{6} absorption and to compare with recent \\ion{O}{6} surveys in the SMC and LMC \\citep{hoopes02,howk02} and in the Galactic halo \\citep{savage00}. They also provide a quantitative estimate of the quantities of neutral and weakly ionized gas in the MB and in the ionized high-velocity cloud HVC $291.5-41.2+80$ \\citep{lehner01b}. ",
        "conclusions": "\\subsection{Molecules and star formation in the Magellanic Bridge}\\label{disc1} \\citet{demers98} found using deep $BV$ CCD photometry that the stars in form of stellar clusters in the MB were formed some 10 to 25 Myr ago. Spectroscopic studies of some of these stars confirm that they are massive young stars but with a metallicity $Z \\approx 0.08 Z_\\odot$ \\citep{rol99}, a factor $\\sim$3 lower than the present day metallicity of the SMC. Their location in the MB also rejects the idea that these stars could be runaway stars from the SMC since they could not travel such large distances over their short life-times. It is generally accepted that the MB was most likely pulled from the wing of the SMC some 200 Myr ago, during a close encounter between the two Clouds, whereas the Magellanic Stream was created by a SMC-LMC-Galaxy close encounter $\\sim1.5$ Gyr ago \\citep{gardiner}. Searches for stars in the Stream have, however, produced largely negative results \\citep[see][and references therein]{irwin90}. Interstellar study toward two sight lines in the Stream indicates an abundance similar to the abundance in the SMC, $Z \\approx 0.25 Z_\\odot$ \\citep{lu98,sembach01,gibson00}. Yet, because the present-day MB abundance is a factor $\\sim$3 lower than the present day metallicity of the SMC, \\citet{rol99} postulated that the MB was formed either much earlier ($\\sim$8.5 Gyr ago) from more primordial gas, or  some 0.2 Gyr ago from a mixture of SMC gas and an unenriched (primordial) component. Toward both the Magellanic Stream and Bridge, molecules were found, yet because star formation occurs in the MB and not in the Stream, the physical conditions must be different. Even though young stellar clusters exist in the MB \\citep{demers98}, their formation is not very efficient and only happens in recent times, since the MB gas is poor in metals. \\citet{christ97} suggested that the MB was formed earlier than the Stream, in order to explain why star formation occurs only in the MB. In such a scenario, the low metallicity is explained by the tidal stripping of gas from SMC in its earlier chemical evolution, and stars form late, only after tidally induced cloud mergers have had a chance to occur in substantial numbers, and to produce some massive clouds. Such scenario is supported by some recent \\ion{H}{1} Parkes All Sky Survey (HIPASS) data of the entire Magellanic Stream which indicate dual filaments that are likely to be relics from gas stripped from the SMC and the MB \\citep{putman02}. These observations imply that at least some part of the Stream could be younger than the MB, and that the MB could be older than usually assumed  \\citep[see,][]{putman02}. Pairing of galaxies are very common \\citep{barnes92}, and the SMC and LMC could have been paired since the beginning of their existence, and therefore some part of the MB, if not all, could be much older than previously thought.   The discovery of molecular hydrogen in the MB implies that it is formed {\\em in situ} and/or was tidally stripped from the SMC. Formation of H$_2$ is thought to occur on interstellar dust grains when hydrogen atoms are adsorbed onto the grain surface and react. An hydrogen molecule is then ejected in the gas phase \\citep[e.g.,][]{hollenbach71}. Therefore the formation mechanism is related to the quantity, type of grains, gas density, and metallicity. Dissocation of  H$_2$ depends on the intensity of the ultraviolet radiation field, the presence of hot gas, and the possible effect of shocks. Grains serve as catalysts for the formation of H$_2$ but also as shields against the ultraviolet radiation. \\citet{tumlinson02} found that for the LMC and SMC low metallicity gas differences in the formation and destruction balance of H$_2$: grain formation rates about 1/3 to 1/10 of the Galactic value and ultraviolet radiation field 10 to 100 times the Galactic mean value. The H$_2$ formation can be written as $t = 1/(Rn_{\\rm H})$ \\citep[e.g.,][]{shull82}, where $R$ is the formation rate coefficient, ($R= 1-3\\times 10^{-17}$ cm$^{3}$\\,s$^{-1}$ for Galactic conditions, $R \\approx 0.3\\times 10^{-17}$ cm$^{3}$\\,s$^{-1}$ for LMC, SMC conditions) and $n_{\\rm H}$ the hydrogen column density. Assuming LMC, SMC conditions for the MB, $t = 11/n_{\\rm H} $ Gyr. Along the DI\\,1388 sight line, $n_{\\rm H}$ cannot be high enough to lower $t$ to the believed 200 Myr year old of the MB, and therefore from those assumptions most of the observed  H$_2$ has survived the tidal stripping; unless if the MB is much older.   Alternatively, H$_2$ could have  (partially) formed in small compact and dense cloudlets within the MB. Formation of clusters of massive stars occurs in the MB \\citep[][and references therein]{rol99,demers98}, and if the formation processes are similar to that observed in the Galaxy, this requires very high density clouds and cloud-cloud collisions \\citep{scoville86,evans99}. In this picture H$_2$ can form rapidly in dense regions and disperse in more diffuse clouds. Therefore the observed diffuse molecular clouds along the DI\\,1388 sight line could trace or be the remnant of denser molecular clouds, the latter being the place where star formation occurred. The detection of cold \\ion{H}{1} cloud with spin temperatures between 20 and 50 K in the MB by \\citet{kob99} also supports the idea that some regions must be at high densities ($\\ga 100$ cm$^{-3}$). The pressure in the MB must be low on average, yet tidal interactions between the SMC and LMC could have created strong enough perturbations in the overall low density MB gas to produce high density peaks, conditions not existing in the Stream because possibly the interaction in the MB involves closer galaxies and/or is over a longer time.  \\subsection{N/O ratio in the Magellanic Bridge}\\label{disc2} The nucleosynthetic origin of N has been subject of large debate, but its understanding is fundamental for comprehending the chemical evolution of galaxies \\citep{vila93,henry99}. Nitrogen is mainly produced in the six steps of the CN branch of the CNO bi-cycle within H burning stellar zones, where $^{12}$C serves as a catalyst. First generation stars produce their own C during the He burning phase, and N production must be fairly independent of the initial composition of the star of which it is synthesized, and the synthesis is said {\\em primary}. Beyond the first generation of stars, the gas from which these  stars formed is already polluted with C and O. The amount of N in this material will be proportional of its C abundance, and the N synthesis is said to be {\\em secondary} in this case. Thus, primary nitrogen is independent of the metalliticity, and secondary nitrogen increases with increasing metallicity. Most measurements of N/O in our Galaxy and nearby Galaxies are made in stars and \\ion{H}{2} regions and only few direct measurements in the gas are available in absorption, and therefore our measurement in the MB is particularly notable.  Toward DI\\,1388, the MB N/O ratio is $[{\\rm N/O}] = -0.05 \\pm 0.02$ (from \\ion{N}{1} and \\ion{O}{1} column densities), i.e. near a solar value. \\citet{lehner01a} suggested that N might be deficient in the MB, however, they also noted that this could be mainly due to photoionization effect. The {\\fuse}\\/ data allow us to have precise measurements of \\ion{N}{1} and \\ion{O}{1} but also give access to ionized N via \\ion{N}{2} and \\ion{N}{3}. If \\ion{N}{1} is preferentially ionized with respect to \\ion{O}{1}, $[{\\rm N/O}]$ would be solar or even above solar. The N/O ratio in the gas is in agreement with the results on the photospheric abundances of the B-type stars analysis where $[{\\rm N/O}] \\approx -0.08$ \\citep[from a differential abundance analysis,][]{rol99}. In contrast, in both the LMC and SMC, N in the \\ion{H}{2} regions is systematically underabundant by a factor 4--6 compared to the main sequence B-type stars \\citep[see,][]{garnett}. An analysis of the \\fuse\\/ data using absorption lines could help to constrain N/O in the ISM of the LMC and SMC, and therefore help to comprehend these differences.  \\citet{henry00} compiled a survey of Galactic and extragalactic \\ion{H}{2} regions and stars to analyse the behavior of N/O as a function of O/H. They observe a plateau (though with some scatter) of N/O measurements when $[{\\rm O/H}] < -1$ suggesting independence of N/O with metallicity, consistent with primary nitrogen formation; and when $[{\\rm O/H}] \\ga -1$,  a rise with a large scatter of N/O with metallicity is observed, more suggestive of secondary nitrogen formation. The MB metallicity \\citep[$-1.1 \\pm 0.1$ dex for C, N, O, Mg, and Si in the photosphere of main sequence B-type stars,][]{rol99} is at this apparent boundary. \\citet{henry00} found, at the MB metallicity, N/O lower by a factor 3 to 10 compared to the MB. Yet, the typical error bars for these measurements are large, typically $\\sim \\pm 0.4$ dex for $\\log ({\\rm N/O})$ and $\\sim \\pm 0.2$ dex for $\\log ({\\rm O/H})$, and within the error bars few measurements of N/O in \\ion{H}{2} regions could have a solar value for a MB metallicity. In the damped Ly$\\alpha$ systems, at  MB metallicity (or slighly higher or much lower than the MB metallicity), N/O is always at least a factor 3 lower than solar \\citep{prochaska02,lu98b}.  From both the studies of the gas and the massive stars, N/O is much higher in the MB than N/O values in the damped Ly$\\alpha$ systems, implying that N in the MB is more secondary than primary. Most of nitrogen is formed in intermediate mass-stars (1--8 M$_\\odot$) with roughly a charateristic lag time of 250 Myr \\citep{henry00}. If the MB is 200 Myr old (see \\S~\\ref{disc1}), most of the nitrogen must have come from an N-enriched region of the SMC. \\citet{henry00} argue that since the distribution of N/O at a single O/H value appears to be clustered toward low N/O values, the high values of N/O might be due to nitrogen enrichment by Wolf-Rayet stars or luminous blue variables."
    },
    "0206/astro-ph0206470_arXiv.txt": {
        "abstract": "{ The X-ray emission of RXJ1856.5-3754 has been found to coincide to unprecedented accuracy with that of a blackbody, of  radius $ 5.8\\pm0.9$ km for the measured parallax distance of $ 140\\,$pc (Burwitz et al. 2001, Drake et al. 2002). If the emission is uniform over the whole surface of a non-rotating star, the mass of the star cannot exceed $0.75\\pm0.12M_\\odot$ regardless of its composition. If the compact object is a quark star described by the MIT-bag equation of state (a ``strange star''), the mass is no more than $0.3M_\\odot$. Comparably small masses are also obtained for the X-ray bursters Aql X-1 and KS1731-260 for some fits to their spectra. ",
        "introduction": "As noted by several authors, conventional neutron stars always have a (circumferential) radius larger than 6 km.  Recent reports of a rather small blackbody radius of a nearby neutron-star candidate have generated speculation that the compact object may not be a neutron star but a quark star instead. Here, we point out that although the actual composition of stars with a 6 km radius is not known, what would make such stars unusual is their low mass, posing a challenge to current theories of their formation. Detailed simulations of supernovae do not predict remnant masses less than $ ~1.2 M_\\odot$ (Timmes et al. 1996). ",
        "conclusions": ""
    },
    "0206/astro-ph0206193_arXiv.txt": {
        "abstract": "To better understand diffuse ionized gas kinematics and halo rotation in spiral galaxies, we have developed a model in which clouds are ejected from the disk and follow ballistic trajectories through the halo.   The behavior of clouds in this model has been investigated thoroughly through a parameter space search and a study of individual cloud orbits. Synthetic velocity profiles have been generated in $z$ (height above the plane) from the models for the purpose of comparing with velocity centroid data from previously obtained long-slit spectra of the edge-on spirals NGC 891 (one slit) and NGC 5775 (two slits).  In each case, a purely ballistic model is insufficient in explaining observed DIG kinematics. In the case of NGC 891, the observed vertical velocity gradient is not as steep as predicted by the model, possibly suggesting a source of coupling between disk and halo rotation or an outwardly directed pressure gradient.  The ballistic model more successfully explains DIG kinematics observed in NGC 5775; however, it cannot explain the observed trend of high-z gas velocities nearly reaching the systemic velocity.  Such behavior can be attributed to either an inwardly directed pressure gradient or a possible tidal interaction with its companion, NGC 5774. In addition, the ballistic model predicts that clouds move radially outward as they cycle through the halo.  The mass and energy fluxes estimated from the model suggest this radially outward gas migration leads to a redistribution of material that may significantly affect the evolution of the ISM. ",
        "introduction": "The majority of the ionized gas in the interstellar medium of the Milky Way resides in a vertically extended layer known as the Reynolds Layer or Warm Ionized Medium (WIM).  Photoionization by massive stars in the disk is likely the dominant source of energy input into the layer given its energetic requirements (Reynolds 1993) as well as the first order agreement between observed line ratios (Haffner, Reynolds, \\& Tufte 1999) and photoionization models (Domg\\\"{o}rgen \\& Mathis 1994; Sokolowski 1994; Bland-Hawthorn et al. 1997; Sembach et al. 2000).  In external galaxies, where the layer is more commonly referred to as diffuse ionized gas (DIG), early work concentrated on the characterization and occurrence of these layers for ``normal'' galaxies (e.g. Rand, Kulkarni, \\& Hester 1990; Dettmar 1992; Ferguson, Wyse, \\& Gallagher 1996). A number of observational results suggest that local levels of star formation are responsible for these layers both in energizing the gas and in accelerating the gas upward via supernova activity and winds. Among the results supporting this conclusion are the correlation between layer prominence (in brightness and extent) and far-infrared luminosity (e.g. Rand 1996; Rossa \\& Dettmar 2000), and with tracers of star formation at other wavelengths (e.g. Dahlem, Dettmar, \\& Hummel 1994). Further evidence linking DIG halos to star formation includes the identification of DIG structures resembling the chimneys described by Norman \\& Ikeuchi (1989), which form when expanding supernova-driven supershells break through the main gaseous layers of a disk (e.g. Collins et. al. 2000). In recent years, the edge-ons NGC 891 and NGC 5775 have provided a wealth of information on the origin and excitation of gaseous halos.  Each of these galaxies has relatively high rates of star formation as indicated by their far-infrared luminosities.  NGC 5775 is classified as a starburst galaxy (Condon \\& Broderick 1988), though there is little evidence for a bright, nuclear starburst.  DIG halos in these galaxies are considerably more extended in $z$ than the Reynolds Layer, and in the case of NGC 5775, DIG can be detected up to $z\\approx13$ kpc (Rand 2000).  In each of these galaxies, line ratio data indicate that a non-negligible fraction of DIG halo emission originates in a component that is either energized by an additional ionizing source such as shocks (e.g. Shull \\& McKee 1979) or is at a considerably higher temperature (e.g. Reynolds, Haffner, \\& Tufte 1999) than the bulk of halo DIG (Rand 1998; T\\\"{u}llman et. al. 2000; Collins \\& Rand 2001).  Though significant research has been carried out regarding their energization, relatively little work has been pursued on the issue of rotation and support of these vertically extended layers through studies of DIG kinematics.  Current evidence suggests that halo gas rotates more slowly than gas in the disk.  In the case of NGC 891, the HI halo component appears to rotate 25-100 km s$^{-1}$ more slowly than the disk component (Swaters et. al. 1997).  In addition, the ``bearded'' HI position-velocity map of the intermediately inclined galaxy NGC 2403 has been successfully modeled as a superposition of a thin disk and a more slowly rotating thick halo (Schaap et. al. 2000; Fraternali et al. 2001). At higher-redshift, halo kinematics has been investigated  through quasar absorption-line studies (e.g. Charlton \\& Churchill 1998).  The recent results of Steidel et al. (2002) are consistent with a kinematic model where the QSO-absorbing gas resides in a halo which rotates more slowly than its parent galaxy.  However, though HI data cubes yield information for gas close to the midplane, they shed little light on gas kinematics at highest $z$ for low-redshift galaxies, which is better traced by the more extended ionized gas layer.  In NGC 891 (Rand 1997) and NGC 5775 (Rand 2000; T\\\"{u}llmann et. al. 2000), emission line centroids indicate a slow-down in DIG rotation with $z$. However, in the case of NGC 891, Benjamin (2000) has found that the observed drop-off in rotation speed from $z=1$ to 5 kpc of $\\sim30$ km s$^{-1}$ is considerably more shallow than the value of 80 km s$^{-1}$ determined by assuming a fluid disk in hydrostatic equilibrium and calculating rotation speeds in a standard galactic potential.  This calculated value is only weakly dependent on the choice of potential and is predominantly due to the geometric effect of a decrease in the projection of the radial gravitation vector with height above the disk.  Higher order effects included in hydrodynamical models of halo kinematics, such as outwardly directed pressure gradients or a coupling between disk and halo through magnetic tension or viscosity, may resolve this discrepancy.  However, in practice these effects are not well constrained given the lack of understanding on issues such as galactic magnetic field structure and gas densities.  Though a fluid disk in hydrostatic equilibrium is perhaps the simplest model one can adopt in considering halo kinematics, it assumes that the observed gas is dynamically coupled to the surrounding gas. This is not necessarily the case.  An alternate point of view is that extraplanar gas consists of density concentrations sufficiently large that the motion is essentially ballistic.  These clouds then cycle through the halo similar to Bregman's (1980) galactic fountain return flow. Among the pieces of evidence that support such a picture is the fact that DIG layers are considerably thicker (by more than an order of magnitude) than the thermal scale height of 10$^{4}$ K gas in the gravitational field (e.g. Dettmar 1992) indicating turbulent motion and outflows, or, at the very least, a source of pressure support to maintain these thick layers.  This fact, coupled with strong evidence of outflows in normal galaxies (e.g. Golla \\& Hummel 1994; Collins et al. 2000), indicate extensive cycling of material from disk to halo.  Recent observations of a quasar projected at a height of $z=5$ kpc behind the halo of NGC 891 indicate a filling factor of low-ionization species (\\ion{Mg}{2}\\ and \\ion{Fe}{2}) gas of around $f\\sim0.025$ (E. Miller, private communication).  This suggests a very clumpy medium if in fact this gas traces the DIG component, and lends some justification to treating the ISM as a collection of circulating clouds.  Towards the goal of a better understanding of DIG kinematics and halo rotation, we have developed an ISM model which assumes DIG is distributed in clouds which take ballistic trajectories as they move from disk to halo and back.  We then attempt to compare velocity centroid data, obtained from previously presented long-slit data for NGC 891 (Rand 1997) and NGC 5775 (Rand 2000), with this ballistic cloud model.  In all cases the slit is oriented perpendicular to the midplane so that these issues can be addressed as a function of $z$. Further details of the observations can be found in the relevant references.  The ballistic model is discussed in \\S\\ 2.  Data are compared to the model and results are discussed in \\S\\ 3. Conclusions are summarized in \\S\\ 4. ",
        "conclusions": "  1. The vertical velocity gradient observed in NGC 891 is not as steep as predicted by the ballistic cloud model.  This suggests the presence of drag between disk and halo such as through magnetic tension or viscous interactions between clouds. Alternatively, an outwardly directed pressure gradient could explain the gas kinematics.  2. The ballistic model is more successful in explaining DIG kinematics in NGC 5775.  The filamentary morphology of some of the DIG emission in this galaxy, suggesting the presence of outflows, may explain why the ballistic model provides a reasonable representation.  However, further simulations of chimneys in a rotating frame would be necessary to confirm this speculation.   The ballistic model completely fails at high-$z$ where velocities nearly reach systemic. An inwardly directed pressure gradient may provide the extra support needed to explain the apparent slow rotation at high-$z$. A possible tidal effect due to its companion galaxy, NGC 5774, should be considered when interpreting the kinematics observed in NGC 5775.  3. The ballistic model predicts that clouds migrate radially outward as they cycle through the halo.  The mass fluxes estimated from the models of NGC 891 and NGC 5775 imply that significant amounts of gas can be involved in these migrations. Such migrations could cause a redistribution of gas that could affect metallicity gradients as well as star formation properties. Such effects have been previously investigated by Charlton \\& Salpeter (1989), for example, but extensive observations of the kinematic behavior of edge-on galaxies should yield important constraints on such redistribution.  4. We are limited in this work by the fact that we have data for only a few slit positions. Hence, we are beginning to carry out Fabry-Perot observations of DIG in edge-on galaxies which will allow us to obtain kinematic information with full two-dimensional coverage.  Such work will allow issues of halo rotation to be addressed more completely."
    },
    "0206/physics0206021_arXiv.txt": {
        "abstract": " ",
        "introduction": "A charge-coupled device (CCD) is widely used in an optical imaging as well as an X-ray imaging. When an X-ray photon is photoabsorbed in the CCD, it generates a primary charge cloud consisting of thousands of electrons. The size of the primary charge cloud is 73$\\times(E/2.3 \\,{\\rm keV})^{1.75}$\\,nm where $E$ is energy of X-ray photon in unit of keV~\\cite{charge_cloud}. The primary charge cloud expands through diffusion process until they reach the potential well of the CCD pixel. The final charge cloud after diffusion process is collected into several pixels forming various types of event pattern depending on how they split.  We have developed a new technique ``mesh experiment'' which enables us to restrict the X-ray point of interaction with a subpixel resolution~\\cite{tsunemi97}.  Hiraga {\\it et al.}~\\cite{hiraga98} investigated the event pattern produced and measured the charge cloud shape produced in the front-illuminated (FI) CCD for the first time. They found that a charge cloud shape could be well represented by a Gaussian function.  They also obtained the standard deviation, $\\sigma$, of the final charge cloud to be $0.7 \\sim 1.5 \\,\\mu$m for $1.5\\sim 4.5$\\,keV X-rays by using a FI CCD.  Based on this experiment, they confirmed that there are three parameters tightly coupled together~\\cite{tsunemi00}; an X-ray point of interaction within a pixel, an X-ray event pattern and a charge cloud shape.  Any two parameters can determine the third one.  The event pattern is quite easily noticed while the charge cloud shape is difficult to be measured.  Currently, they can separately measure it only by using the mesh experiment.  Therefore, they can determine the X-ray point of interaction with much better spatial resolution than the pixel size. They obtained the position resolution of 0.7\\,$\\mu$m using the pixel size of 12\\,$\\mu$m if an X-ray photon became a split pixel event. However, due to its relatively small size of the charge cloud, the fraction of split pixel events whose point of interaction can be improved is less than 10\\,\\% of the total events.  In this paper, we applied the mesh experiment to the back-illuminated (BI) CCD.  Figure~\\ref{fig:fi_cross_section} shows the cross-section and the potential profile of FI CCDs. In the case of FI CCDs, since X-ray photons enter from the front side where there are electrodes.  A major part of X-rays are photo-absorbed close to electrodes, resulting a final charge cloud to be relatively small. Figure~\\ref{fig:bi_cross_section} shows the cross-section and the potential profile of BI CCDs.  In the case of BI CCDs, a substrate of CCD, shown in Fig~\\ref{fig:fi_cross_section} (a), is etched from the back toward the depletion layer (thinning process) in order to achieve high quantum efficiency for blue lights and low energy X-rays. After the thinning process, almost all of the substrate has been removed. The thickness of the BI CCD employed is $\\simeq20\\,\\mu$m.  Since X-ray photons enter from the back surface of the CCD shown in Fig~\\ref{fig:bi_cross_section}, a primary charge cloud is formed far from the electrodes.  The primary charge cloud expands through diffusion process until they reach the potential well which is just below the electrodes. Therefore, the diffusion time for the charge cloud generated by X-rays is longer than those in the FI CCD, resulting that a larger charge cloud shape is expected.  Low energy X-rays, however, are photo-absorbed in the field-free region of BI CCDs. The primary charge cloud expands due to a diffusion process toward the electrodes as well as the back surface. The charge traveled to the back surface will be lost due to a recombination process. In the production of BI CCDs, accumulation is the most significant process that gives an internal potential to the back surface of CCDs and repels the signal charge to the electrodes~\\cite{accumulation1}.  In the case of the BI CCD employed, the accumulation process is performed by ion implantation~\\cite{accumulation2}.  As shown in Fig~\\ref{fig:bi_cross_section} (b), the potential at close to the back surface is locally lower than that at the inner region caused by the accumulation process which enables us to collect the charge generated close the back surface. ",
        "conclusions": "We have performed the mesh experiment with the BI CCD and measured the charge cloud size for the first time. We found that there are two components of the charge cloud shape to reproduce the data obtained all for Mo-L, Al-K and Ti X-rays: a narrow component and a broad component. The narrow component possesses a $\\sigma$ of $2.8 - 5.7\\, \\mu$m which strongly depends on the attenuation lengths of incident X-ray energies. The broad component, on the other hand, possesses much larger $\\sigma$ than that of the narrow component. It is $\\simeq 13\\,\\mu$m and almost independent of the attenuation length of X-rays. The accumulation technique to fabricate BI CCDs is not employed in the FI CCDs and the charge produced even in the field-free region can be collected toward electrodes. We thus suppose that the narrow and the broad components are originated in the depletion region and the field-free region, respectively. This hypothesis is consistent with both by the design value of the CCD and by the fraction of X-ray events for both components.  We have found that the linearity between the incident X-ray energies and the pulse height is not good. The pulse height of Mo-L X-rays is $\\simeq 30$\\,\\% lower than that expected.  This might be explained by the fact that the effect of the recombination with the impurities is significant in the accumulation region since the attenuation length of Mo-L is shortest among those X-rays employed.  \\acknowledgement  J.H. is partially supported by JSPS Research Fellowship for Young Scientists, Japan. This work is partly supported by the Grant-in-Aid for Scientific Research by the Ministry of Education, Culture, Sports, Science and Technology of Japan (13874032, 13440062)."
    },
    "0206/astro-ph0206079_arXiv.txt": {
        "abstract": "Quasar-driven winds are currently the best candidates for accounting for the pre-heating of the intergalactic medium in clusters. Such winds, occurring during early phases of the evolution of spheroidal galaxies, shock-heat the interstellar gas, thus inducing a detectable Sunyaev-Zeldovich effect. We estimate the amplitude and the angular scale of such effect as well as its counts as a function of the comptonization parameter $y$. The contamination due to radio emission by the quasar itself is also discussed. The corresponding mean Compton distortion of the cosmic microwave background spectrum is found to be well below the COBE/FIRAS upper limit. ",
        "introduction": "\\label{intro}  It is natural to expect that extremely powerful sources, such as quasars, strongly affect the surrounding medium. % The recent evidences of a tight relationship between black hole masses and velocity dispersions of the host galactic bulges (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000a,b; Merritt \\& Ferrarese 2001a,b; McLure \\& Dunlop 2001) have strongly strengthened the case for a close connection between the evolutionary pathways of spheroidal galaxies and of quasars. Super-massive black holes have been found to be ubiquitous at the centers of local spheroidal galaxies (Kormendy \\& Richstone 1995; Magorrian et al. 1998; van der Marel 1999). The observed correlation between the mass of the black hole and that of the galaxy spheroidal component hints at a substantial feedback of the nuclear activity to the surrounding medium (Silk \\& Rees 1998; Monaco et al. 2000; Granato et al. 2001): energetic quasar-driven winds can sweep up the interstellar gas and halt both star formation and the growth of the central black hole.  If a significant fraction of the huge amount of energy released by quasars % goes into ionization and heating of the neighboring gas, Compton cooling may produce a detectable Sunyaev-Zeldovich (1972; SZ) effect. The idea of strong shock heating of the medium originated by energetic outflows from early quasars, originally developed by Ikeuchi (1981), was recently revived by Natarajan, Sigurdsson \\& Silk (1998), Natarajan \\& Sigurdsson (1999), and Aghanim, Balland \\& Silk (2000). Evidences of strong quasar-driven winds can be seen in Broad Absorption Lines (BAL) quasars, comprising 10\\%--15\\% of optically selected quasars (Hamann \\& Ferland 1999 and reference therein). The dynamic interaction of such energetic outflows with the surrounding proto-galactic gas can heat it to high temperature. A release of an amount of mechanical energy from Active Galactic Nuclei several times larger than that produced by supernovae in the surrounding galaxies may be necessary to account for the pre-heating of the intergalactic medium in clusters (Wu et al. 2000; Kravtsov \\& Yepes 2000; Balogh et al. 2001; Bower et al. 2001). Natarajan \\& Sigurdsson (1999) suggested that the SZ effect associated to very energetic, quasar driven, winds may account for the reported isolated Cosmic Microwave Background (CMB) temperature decrements in directions where no clusters of galaxies are detected, but quasar pairs are present (Jones et al. 1997; Richards et al. 1997). Detection of this effect would obviously be informative on the physics of quasar/galaxy evolution.  In this paper we present a new investigation of the problem, taking explicitly into account the relevant energetics as well as the observed relationships between black hole and host galaxy properties. Our approach, focussing on energetics, is less liable to uncertainties ensuing from the poor knowledge of the details of the gas heating process. % In Section 2 we introduce the basic ingredients entering our calculations and estimate the amplitude and the angular scale of the SZ effect. In Section 3 we present tentative estimates of its number counts. The results are discussed in Section 4. ",
        "conclusions": "\\label{obs}  \\subsection{Effect of local radio emission}  It is entirely plausible, and perhaps required by evidences of a substantial preheating of the intra-cluster gas, that quasars inject in the surrounding medium an amount of energy sufficient to produce a detectable SZ effect. The corresponding angular scale is expected to be relatively small (sub-arcmin), so that the blurring effect by the emission from the quasar itself or from the host galaxy may be important.  To estimate to what extent the radio emission associated with the quasar might blur the SZ signal, we refer to the median ratio of monochromatic luminosities at 5 GHz and at $\\nu_B \\simeq 6.82 \\times 10^{14}\\,$Hz (corresponding to $\\lambda_B = 0.44\\,\\mu$m), for the radio-loud and radio-quiet quasars in Table~2 of Elvis et al. (1994). We find $\\log(l_{5GHz}/l_B)_{\\rm median} = -0.47$ for radio-quiet quasars and $=2.2$ for radio loud.  Taking into account that the contribution to the antenna temperature at the frequency $\\nu$ within a solid angle $\\omega_{SZ} = \\pi \\theta_{SZ}^2$, of a source of 5 GHz flux $S_{5GHz}$ and spectral index $\\alpha\\simeq 0.7$ ($S_\\nu \\propto \\nu^{-\\alpha}$) is: \\begin{equation} \\Delta T_A = {S_{5GHz} (\\nu/5\\,{\\rm GHz})^{-\\alpha} c^2\\over 2 k_b \\nu^2 \\omega_{SZ}} \\ , \\label{DeltaT_A} \\end{equation} we have, for radio-quiet quasars, in the Rayleigh-Jeans region \\begin{eqnarray} {\\Delta T_A \\over T_A} &\\simeq& 1.1 \\, 10^{-3}\\left({E_{\\rm tot}\\over 10^{62}} \\right)^{7/15}\\left({\\epsilon_R\\over 0.1}\\right)^{8/15}h_{50}^{2}\\cdot \\nonumber \\\\ & &\\cdot {10^7 {\\rm yr} \\over t_q} \\!\\! \\left({5\\,{\\rm GHz} \\over \\nu}\\right)^{2+\\alpha} \\left(1+z\\over 3.5\\right)^{-\\alpha}\\ , \\label{TA} \\end{eqnarray} In the case of radio loud quasars the coefficient would be about 500 times higher and therefore the radio emission would easily overwhelm the SZ signal, except, perhaps, at high frequencies ($\\simeq 100\\,$GHz). However, only a minor fraction ($\\lsim 10\\%$) of quasars are radio loud. But even in the case of radio quiet quasars the radio emission may, at least partially, fill up the SZ dip at $\\nu \\lsim 10\\,$GHz, particularly if the quasar is surrounded by an intense starburst, which is also radio bright, as is frequently the case at high $z$ (Omont et al. 2001). The contamination by radio emission sinks down rapidly with increasing frequency, but at $\\nu \\gsim 150\\,$GHz, dust emission powered by star formation in the host galaxy, may take over, due to its spectrum steeply rising with increasing frequency.  On the other hand, the quasar lifetime and the duration of intense starbursts are generally shorter than the time $t_g$ for the blast-wave to reach the boundary of the host galaxy. Therefore the SZ signals may be observed when the quasars that originated them are dead.  \\subsection{Global comptonization distortion}  The mean distortion of the CMB spectrum due to comptonization by the electrons heated by the blast waves, measured by the mean value $\\langle y\\rangle$ of the parameter $y$ integrated along the line of sight, is $1/4$ of the fractional amount of energy injected in the CMB per unit volume: \\begin{equation} \\langle y\\rangle = {f_{\\rm h, eff}\\over 4} \\int \\int_0^\\infty dz\\, d\\log L_B {\\phi_B(L_B,z) k_B L_B \\over \\epsilon_{\\rm CMB}(z)}{t_{\\rm exp}\\over t_c}{dt\\over dz} \\ . \\label{<y>} \\end{equation} where $f_{\\rm h, eff}$ is the effective fraction of the total energy released by the quasars carried by blast waves. From the above discussion it follows that pre-heating of the ICM requires $f_{\\rm h, eff}\\lsim 0.1$. Using the double-power-law quasar luminosity function by Pei (1995), we find: \\begin{equation} \\langle y\\rangle \\simeq 2.4 \\times 10^{-6} {f_{\\rm h, eff}\\over 0.1} \\ . \\label{<y>max} \\end{equation} Using Pei's (1995) exponential luminosity function, the value of $\\langle y\\rangle$ increases by 20\\%. On the other hand, it may be noted that Eq.~(\\ref{<y>}) includes the contribution (which turns out to be quite substantial) of low-$z$ quasars to the integral (the contribution from the redshift range 0--2 is comparable to that from $2\\le z \\le 4$); however, in the present framework, the most energetic blast-waves should be associated to early phases of the quasar/galaxy evolution. Thus the expected global comptonization distortion of the CMB spectrum induced by quasar driven blast waves is well below the COBE/FIRAS limit $\\langle y\\rangle < 1.5 \\times 10^{-5}$ (Fixsen et al. 1996). This limit could be (marginally) exceeded only in the unrealistic case that the total energy emitted by all quasars over the entire life of the universe went into heating of the surrounding gas. This would also entail an excessive (by almost one order of magnitude) pre-heating of the ICM.   \\bigskip\\noindent {\\bf ACKNOWLEDGMENTS}  Work supported in part by MIUR and ASI."
    },
    "0206/astro-ph0206286_arXiv.txt": {
        "abstract": "{ In a previous study (Maurin et al., 2001), we explored the set of parameters describing diffusive propagation of cosmic rays (galactic convection, reacceleration, halo thickness, spectral index and normalization of the diffusion coefficient), and we identified those giving a good fit to the measured B/C ratio. This study is now extended to  take into account a sixth free parameter, namely the spectral index of sources. We use an updated version of our code where the reacceleration term comes from standard minimal reacceleration models. The goal of this paper is to present a general view of the evolution of the goodness of fit to B/C data with the propagation parameters. In particular, we find that, unlike the well accepted picture, and in accordance with our previous study, a Kolmogorov-like power spectrum for diffusion is strongly disfavored. Rather, the $\\chi^2$ analysis points towards $\\delta\\gtrsim 0.7$ along with source spectra index~$\\lesssim 2.0$. Two distinct energy dependences are used for the source spectra: the usual power-law in rigidity and a law modified at low energy, the second choice being only slightly preferred. We also show that the results are not much affected by a different choice for the diffusion scheme. Finally, we compare our findings to recent works, using other propagation models. This study will be further refined in a companion paper, focusing on the fluxes of cosmic ray nuclei. ",
        "introduction": "Cosmic rays detected on Earth with kinetic energies per nucleon from 100 MeV/nuc to 100 GeV/nuc were most probably produced by the acceleration of a low energy galactic population of nuclei, followed by diffusion in the turbulent magnetic field. The acceleration process and the diffusion process have a magnetic origin, so that they should depend on rigidity. The rigidity dependence of the diffusion coefficient is given by quasi-linear theory as \\begin{equation} K({\\cal R}) = K_0 \\beta  \\left( \\frac{\\cal R}{\\mbox{1 GV}} \\right)^\\delta \\end{equation} where the parameters $K_0$ and $\\delta$ should ideally be given by the small-scale structure of the magnetic field responsible for the diffusion.  As this structure is not well observed, some theoretical assumptions must be made in order to predict $\\delta$. As regards the spectrum just after acceleration, the situation is far from clear, as it depends on the details of the acceleration process. Several models give a power-law distribution ({\\em e.g.} Berezhko et al., 1994, Gieseler et al., 2000) \\begin{equation} \\frac{dQ}{dp} \\propto {\\cal R}^{-\\alpha} \\end{equation} with a definite value for $\\alpha$ which depends on the model.  Most analyses of cosmic ray nuclei data assume given power-laws for the diffusion and acceleration energy dependence , so that the results partially reflect certain theoretical {\\em a priori}. In this work, we try to avoid this bias by determining the quantities $\\alpha$ and $\\delta$ directly from the data, in particular B/C, for reasons exposed below.  The paper is organized as follows. We first recall the main features of our diffusion model. As a few modifications have been made since previous works, $\\S$\\ref{niou} is devoted to their description and justification. Then, the analysis method is described in $\\S$\\ref{reunnes} and the results are shown and discussed in $\\S$\\ref{reseultses}; a comparison is eventually made with other similar works in $\\S$\\ref{compar}. ",
        "conclusions": "Forgetting for a while some of our theoretical {\\em a priori} about the diffusion power spectrum, a new picture of cosmic ray propagation seems to emerge, motivated by the B/C analysis. In this new picture, high values for the diffusion coefficient spectral index ($\\delta\\gtrsim 0.6-0.7$) and source spectral indices $\\alpha \\sim 2.0$ are favored. This latter result is rather satisfactory: as emphasized in a recent working group report on SNR shocks (Drury et al., 2001), even {\\em if nonlinear acceleration models do not produce precise power-law spectra [...] the effective differential energy spectral index is close to 2.0. Furthermore, as pointed out in a series of papers by Vainio and Schlickeiser (2001, and see references therein), diffusive shock wave acceleration naturally yields smaller values of $\\alpha $ if the correct scattering center compression ratio is used instead of the gas compression ratio.\"\\\\} This trend should be carefully analysed and discussed in the light of measured differential fluxes, in order to confirm or point out the possible inconsistencies in the current propagation treatments (see companion paper, Donato, Maurin \\& Taillet, in preparation). Briefly, the major arguments against large $\\delta$ come from anisotropy measurements at high energy and from theoretical preference for Kolmogorov-like turbulence spectra. However, Ptuskin et al. (1997) -- in their self-consistent analytical propagation model including gas, cosmic ray and magnetic field -- derived $\\delta\\sim 0.55$, $\\alpha\\sim 2.1$ and argue that the observed anisotropy could be as well due to a particularity of the local structure of the Galactic magnetic field. Theoretical objections against too high values of $\\delta$ are probably more robust.   For the rest, the conclusions of this paper can be summarized as follows : (i) we performed for the first time a full analysis of diffusion/convection/reacceleration models in the whole 6-dimensional parameter space ($\\alpha$, $\\delta$, $K_0$, $L$, $V_a$, $V_c$), and $\\delta\\sim 0.7-0.9$; the values $\\alpha\\sim2.0$ are preferred; (ii) this preference holds whatever the specific form of the spectrum at low energy; the numerical values of the other parameters are also only slightly modified by this low energy dependence even though deviation from a power-law at low energy is preferred. The study of fluxes should give a more definite answer; (iii) $K_0$ scales logarithmically with $\\delta$ and models with small halos tend to one-dimensional models with a simple relation between $\\mu$, $K_0$, $L$ and $V_c$ (see also Taillet \\& Maurin, in preparation); (iv) several existing models are compared and the qualitative and quantitative differences between them are studied and partially explained."
    },
    "0206/astro-ph0206415_arXiv.txt": {
        "abstract": "{We present results from optical, X-ray and radio observations of two X-ray bright ($L_X \\sim 10^{45}$ \\ergs) galaxy clusters. \\aco\\ is at redshift $z=0.1989$ and has line-of-sight velocity dispersion $\\sigma_v = 1330$ \\kms\\ as measured from 57 cluster galaxies. It has regular X-ray emission without signs of substructure, a Gaussian velocity distribution, lack of a cooling flow region and significant deviations from the observed scaling laws between luminosity, temperature and velocity dispersion, indicating a possible merging shock.  There is only one spectroscopically confirmed cluster radio galaxy, which is close to the X-ray peak. \\\\ \\mbox{1RXS~J131423.6$-$251521} (for short \\rxj) has $z=0.2474$ and $\\sigma_v = 1100$ \\kms\\ from 37 galaxies. There are two distinct galaxy groups with a projected separation of $\\approx 700$ kpc. The velocity histogram is bi-modal with a redshift-space separation of $\\sim 1700$ \\kms, and the X-ray emission is double peaked.  Although there are no spectroscopically confirmed cluster radio galaxies, we have identified a plausible relic source candidate. ",
        "introduction": "Clusters of galaxies form a representative population which traces the highest initial density fluctuation peaks. They are excellent tools for exploring the distant Universe and are used to constrain cosmological models.  A significant fraction of clusters shows evidence of substructure (e.g. Geller \\& Beers \\cite{gb82}, Dressler \\& Shectman \\cite{ds88}, West \\cite{w94}; for a recent review see Pierre \\& Starck \\cite{ps98}) and complexity in the distribution of the different constituents --- galaxies, gas, dark matter (Baier et al.  \\cite{bai96}).  Combined multi-wavelength observations are needed to disentangle the dynamical state of clusters.  Analysis of the velocity and space distribution of the galaxies is very important but, in contrast to the optical, the X-ray analysis is less prone to projection effects and probes better the cluster mass distribution (because the X-ray surface brightness depends on the square of the matter density). The presence of substructure is also revealed in the radio properties of clusters (R\\\"ottgering et al.  \\cite{ro94}, Reid et al.  \\cite{rei98}), but the radio sources in turn may also influence the X-ray emission (Rizza et al.  \\cite{riz00}).  This paper is the third in a series of papers (Pierre et al. \\cite{mp97}, L\\'emonon et al. \\cite{ludo97}) dedicated to studies of distant, bright X-ray clusters discovered in the ROSAT All-Sky Survey (RASS, Voges et al. \\cite{rass}). A sample of $\\approx 10$ clusters with $L_X > 10^{44}$ \\ergs\\ was selected (Pierre et al. \\cite{mp94}) in the redshift range $0.1 < z < 0.3$.  In this paper we present multi-wavelength observations of two of these clusters --- \\aco\\ and \\mbox{1RXS~J131423.6$-$251521} (hereafter shortened to \\rxj).  General data associated with both clusters are given in Table~\\ref{tab1}.  \\begin{table}  \\caption{ Properties of the two clusters.  References: Abell (\\cite{abell}), Abell et al. (\\cite{aco}) for the coordinates and richness; Bautz \\& Morgan (\\cite{bm70}) for classification for \\aco ; coordinates for \\rxj\\ are from RASS (Voges et al. \\cite{rass}); $T_X$ and $L_X$ in the 2--10 keV band are from Matsumoto et al. (\\cite{asca}); redshift and $\\sigma_v$ are from this paper;  the apparent magnitude of an $L^{*}$ galaxy was calculated using $M^{*}(B_j)=-21.8$, obtained by adjusting the Lumsden et al. (\\cite{lum97}) value of $M^{*}(B_j)=-20.16$ to the cosmological parameters used in this paper.}  \\label{tab1} {\\small \\begin{tabular}{lcc} \\hline \\hline & \\aco\\ & \\rxj\\ \\\\ \\hline R.A. (J2000) & 12:03:16.0 & 13:14:23.6 \\\\ Dec. (J2000) & $-$21:30:42 & $-$25:15:21 \\\\ BM Class & III & --- \\\\ Richness & 3 & --- \\\\ Redshift & 0.1989 & 0.2474 \\\\ $\\sigma_v$ km~s$^{-1}$ & 1330 & 1100 \\\\ $B_j^{*}$ & 19.5 & 20.2 \\\\ $L_X\\ (10^{45})$ erg~s$^{-1}$ & 1.5 & 1.8 \\\\ $T_X$ keV & 13.4 & 8.7 \\\\ \\hline \\end{tabular} } \\end{table}  The plan of the paper is as follows: in Section~\\ref{sec:opt} we present optical observations, data reduction, redshift catalogues and data analysis for the two clusters. In Section~\\ref{sec:x} and \\ref{sec:radio} we present ROSAT-HRI observations and data analysis and ATCA radio observations and data analysis respectively.  Finally in Section~\\ref{sec:disc} we discuss the multi-wavelength view of \\aco\\ and \\rxj. Throughout the paper we use $H_0 = 50$ \\kms\\  Mpc$^{-1}$ and $q_0 = 0.5$. ",
        "conclusions": "\\label{sec:disc}  \\subsection{\\aco}  The dynamical state of \\aco\\ is very similar to that of Abell 665 (Gom\\`ez et al.\\ \\cite{g00}) and to the more distant Abell 1300 (L\\'emonon et al.\\ \\cite{ludo97}, Reid et al. \\cite{rei98}), suggesting that it may also be in the final stage of establishing equilibrium after a merger event.  Support for the merger scenario comes from different morphological and physical reasons which are summarized below:  \\begin{itemize} \\item There is no single dominant galaxy. The brightest cluster galaxy (\\#34, Table~\\ref{tab3}) is $35\\arcsec$ away from the X-ray centroid, and 580 \\kms\\ from the cluster mean redshift.\\\\ The putative identification for the central radio source (\\#40, Table~\\ref{tab3}) is offset by 1300 \\kms\\ from the cluster mean redshift. \\item Nearly regular X-ray emission, without substructures but slightly twisted inner part. \\item No cooling flow region. \\item Deviations from the observed scaling relations $L_X$--$T$ and $\\sigma_v$--$T$ (Xue \\& Wu \\cite{xue00}), shown in Table~\\ref{tab:scaling}. The cluster is significantly less luminous than expected from its measured temperature. The very high observed temperature (13.4 keV) is probably an indication of a shock that occurred in the recent past. We cannot exclude a possible overestimation of the published ASCA temperature due to the presence of the background X-ray point source. However, the small difference (in the wrong sense) between the ASCA luminosity and our estimate with the QSO emission excluded (see Table~\\ref{tab:x}) seems unlikely to account for a temperature overestimate of more than 4 keV. A new and more accurate measurement of the temperature is clearly needed. \\item   Based on numerical simulations (Roettiger et al.\\ \\cite{roe97}, Belsole et al.\\ \\cite{bel02}), the regular X-ray morphology and high X-ray temperature are fully compatible with expectations from a past merger. \\end{itemize}  We can also determine the dynamical status of the cluster by comparing its kinetic and potential energies.  From the measured velocity dispersion we find $\\beta_{\\rm spec} = \\mu m_p \\sigma_v^2/kT = 0.84\\pm0.25$, while from the X-ray emission, with the correction factor from Bahcall \\& Lubin (\\cite{bah94}), we have $\\beta^c_X = 1.25 \\beta_X = 0.63\\pm0.1$.  These values are consistent within the uncertainties, indicating that the gas and galaxy motions are close to equipartition.  If a merger occurred recently we might expect signatures at radio wavelengths, such as radio halo/relic sources, and possibly tailed sources (e.g.  En{\\ss}lin et al.\\ \\cite{ens98}, Reid et al.\\ \\cite{rei98}, R\\\"ottgering et al.\\ \\cite{ro94}). There is no evidence for a radio halo, although there is a tailed source (R7) near the cluster centre which could have disrupted it (Giovannini \\cite{gio99}, Liang et al.\\ \\cite{lia00}).  \\begin{table} \\caption{Scaling relations for \\aco\\ and \\rxj. $\\sigma_v$--$T$ is from Xue \\& Wu (\\cite{xue00}), while $L_{\\rm bol}$--$T$ is from Arnaud \\& Evrard (\\cite{arn99}). The temperature from ASCA is in keV, $L_{\\rm bol}$ in units of $10^{44}$ \\ergs\\ and $\\sigma_v$ in \\kms\\ are from this paper.} \\label{tab:scaling} {\\footnotesize \\begin{tabular}{lrcccc} \\hline \\hline Cluster & $T$~ & $\\sigma_v$ & $\\sigma_v$ & $L_{\\rm bol}$ & $L_{\\rm bol}$ \\\\ & (obs) & (obs)      & ($\\sigma_v$--$T$) & (obs)      & ($L_{\\rm bol}$--$T$)     \\\\ \\hline \\\\[-2mm] \\aco    & 13.4 & 1330 & 1670 & 39.8 & 116  \\\\ \\rxj    & 8.7  & 1100 & 1261 & 34.0 & 33.5 \\\\ \\hline \\end{tabular} } \\end{table}   \\subsection{\\rxj}  \\rxj\\ is morphologically and dynamically very different from \\aco.  It shows a clear bi-modal structure --- there are two groups in velocity space separated by $\\sim 1700$ \\kms\\ (cf. Table \\ref{tab6} and Fig.~\\ref{fig5}) which are also separated in the projected galaxy distribution (cf. Fig.~\\ref{fig6}). The dominant galaxies of each group are separated by $\\sim 1000$ \\kms\\ in redshift space, and $2\\arcmin 25\\arcsec$, or $\\approx$700 kpc, in projected distance.  The X-ray emission is elongated, with the centroid located between the two dominant galaxies. The elongation, however, is rotated by $\\approx$20$^{\\circ}$ from the axis connecting the two BCGs.  This may simply be due to the decoupling between the galaxies and gas during the merger.  There are no cluster radio sources within the X-ray extension, with the possible exception of the weak (uncatalogued) source at the position of galaxy \\#48.  If we are witnessing an interaction between two sub-clusters, we might expect stronger radio activity than observed.  However, residual sidelobes from a strong background source $\\sim 7\\arcmin$ south of the centre hamper the detection of any very extended emission. In addition, a more compact ATCA antenna configuration is needed to improve sensitivity to low surface brightness emission.  There are, however, two extended radio sources, one of which (R2) has a steep radio spectrum and no optical counterpart, and is therefore a plausible candidate for a relic source.  The observed $L_X$, $T$ and $\\sigma_v$ for \\rxj\\ are in good agreement with the $L_X$--$T$ and $\\sigma_v$--$T$ scaling relationships (Table~\\ref{tab:scaling}), suggesting that the merger has progressed to the stage where the transient shock heating and radio activity have dissipated.  On the other hand, if the cluster is in a pre-merging phase, then it is unusual that the X-ray elongation is not aligned with the group centres and that there is no sign of X-ray substructure around the eastern group, as revealed, for example, in numerical simulations (Roettiger et al. \\cite{roe97}, Takizawa \\cite{tak00}). The scattered appearance of the projected galaxy distribution of the eastern group compared to the western group (Fig.~\\ref{fig6}) also supports a post-merging scenario.  \\bigskip  In conclusion, our observations suggest that neither cluster is relaxed following a recent merger. However, their properties and scaling laws are quite different, illustrating the diversity in the merging and relaxation processes in cluster formation and evolution.  The current data for the two clusters are compatible with the expectations from the merger of a small group with a bigger cluster for \\aco, and nearly equal mass groups for \\rxj.  Deep XMM and Chandra observations, coupled with detailed numerical simulations are needed to assess these hypotheses and better understand the many aspects of the physical processes occurring during accretion and relaxation over the course of a cluster merger."
    },
    "0206/astro-ph0206309_arXiv.txt": {
        "abstract": "{ We present new near-infrared ($JHK$) bispectrum speckle-interferometry monitoring of the carbon star \\object{IRC+10216} obtained between 1999 and 2001 with the SAO 6\\,m telescope. The  $J$-, $H$-, and $K$-band resolutions are 50\\,mas, 56\\,mas, and 73\\,mas, respectively. The total sequence of $K$-band observations covers now 8 epochs from 1995 to 2001 and shows the dynamic evolution of the inner dust shell. The present observations show that the appearance of the dust shell has considerably changed compared to the epochs of 1995 to 1998. Four main components within a 0\\farcs2 radius can be identified in the $K$-band images. The apparent separation of the two initially brightest components A and B increased from $\\sim 191$ mas in 1995 to $\\sim 351$ mas in 2001. Simultaneously, component B has been fading and almost disappeared in 2000 whereas the initially faint components C and D became brighter (relative to peak intensity). The changes of the images can be related to changes of the optical depth caused, for instance, by mass-loss variations or new dust condensation in the wind. Our recent two-dimensional radiative transfer model of \\object{IRC\\,+10\\,216} suggests that the observed relative motion of components A and B is not consistent with the outflow of gas and dust at the well-known terminal wind velocity of 15 km\\,s$^{-1}$. The apparent motion with a deprojected velocity of 19 km\\,s$^{-1}$ on average and of recently 27\\,km\\,s$^{-1}$ appears to be caused by a displacement of the dust density peak due to dust evaporation in the optically thicker and hotter environment. The present monitoring, covering more than  3 pulsation periods, shows that the structural variations are not related to the stellar pulsation cycle in a simple way. This is consistent with the predictions of hydrodynamical models that enhanced dust formation takes place on a timescale of several pulsation periods. The timescale of the fading of component B can well be explained by the formation of new dust in the circumstellar envelope. ",
        "introduction": "\\label{Sintro} The carbon star \\object{IRC\\,+10216} is a long-period variable star evolving along the Asymptotic Giant Branch (AGB). It is the nearest carbon star known (distance $\\sim 110-150$\\,pc; Crosas \\& Menten \\cite{CroMen97}, Groenewegen et al.\\ \\cite{GroeEtal98}) and the brightest 12\\,$\\mu$m object outside the solar system (IRAS 1986). A strong stellar wind has led to an  almost complete obscuration of the star by dust. The mass-loss rate as measured in CO rotational lines amounts to $2-5 \\times 10^{-5} M_{\\odot}$/yr (Loup et al. \\cite{LoupEtal93}). Detailed two-dimensional (2D) radiative transfer modeling shows that \\object{IRC\\,+10216} had recently suffered from an even higher mass-loss rate of $\\sim 10^{-4} M_{\\odot}$/yr (Men'shchikov et al. \\cite{mbbow2001}). Based on the high mass-loss rate, long period of $P=649$\\,d (Le Bertre \\cite{LeB92}), and carbon-rich  chemistry of the dust-shell, \\object{IRC\\,+10\\,216} is obviously in a very advanced stage of its AGB evolution (see, e.g., Bl\\\"ocker \\cite{Bloe99}).  \\begin{table*} \\caption{Observational parameters. JD refers to the Julian date and $\\Phi$ to the photometric phase. Phase 0 (maximum light) corresponds to JD=2449430 (Mar 18, 1994) as extrapolated from the light curve of Le Bertre (\\cite{LeB92}) with $P=649$\\,d. $\\lambda_\\mathrm{c}$ is the central wavelength and  $\\Delta\\lambda$ the FWHM bandwidth of the filters. $N_\\mathrm{T}$ and $N_\\mathrm{R}$ are the numbers of \\object{IRC +10 216} speckle interferograms and reference-star speckle interferograms, respectively. $T$ is the exposure time per frame, $S$ is the seeing (FWHM), $p$ is the pixel size, and $R$ is the resolution. In the last column, the reference stars are given. The observations of 1995--1998 were already presented in Paper II % and are given here for completeness. } \\label{obstab} \\begin{center} \\begin{tabular}{ % l% c% c% c% c% c% r% r% r% c% c% c l } \\hline & Date & JD & $\\Phi$ & $\\lambda_\\mathrm{c}$ & $\\Delta\\lambda$ & $N_{\\rm T}$ & $N_{\\rm R}$ & $T$\\hspace*{1.2mm}& $S$ & $p$ & $R$ & Ref.\\ star\\\\ &      &    &      & [$\\mu$m]             &  [$\\mu$m]         & &             & [ms]  & [$^{\\prime\\prime}$]  & [mas] & [mas] & \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} $J$ & ~2 Apr 1996 & 2450176 & 1.15 & 1.24 & 0.28 & 1196 &  981 & 200  & 1.2 & 14.6 & 149 & \\object{HIP 51133} \\\\ & 10 Mar 2001 & 2451979 & 3.93 & 1.24 & 0.14 & 1042 &  783 & 160  & 1.0 & 13.3 & 50  & \\object{HD 83871} \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} $H$ & 23 Jan 1997 & 2450472 & 1.61 & 1.64 & 0.31 & 1665 & 2110 & 100  & 1.5 & 19.8 & 70  & \\object{HIP 52689} \\\\ & 10 Mar 2001 & 2451979 & 3.93 & 1.65 & 0.32 &  607 &  915 &  30  & 1.0 & 20.1 & 56  & \\object{HD 83871} \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} $K$ & ~8 Oct 1995 & 2449999 & 0.88 & 2.12 & 0.02 &  251 &  266 & 100  & 1.5 & 31.5 & 92  & \\object{SAO 116569} \\\\ & ~3 Apr 1996 & 2450177 & 1.15 & 2.17 & 0.33 & 1403 & 1363 &  70  & 2.5 & 14.6 & 82  & \\object{HIP 51133} \\\\ & 23 Jan 1997 & 2450472 & 1.61 & 2.19 & 0.41 & 2165 & 1539 &  50  & 0.9 & 30.6 & 87  & \\object{HIP 52689} \\\\ & 14 Jun 1998 & 2450979 & 2.39 & 2.17 & 0.33 &  800 &  571 &  50  & 1.6 & 30.6 & 87  & \\object{HIP 50792} \\\\ & ~3 Nov 1998 & 2451121 & 2.61 & 2.20 & 0.20 & 1087 &  842 &  40  & 1.3 & 27.2 & 75  & \\object{HIP 49583} \\\\ & 24 Sep 1999 & 2451446 & 3.11 & 2.12 & 0.21 & 2702 & 1383 &  80  & 0.9 & 26.4 & 73  & \\object{HIP 49637} \\\\ & 15 Oct 2000 & 2451833 & 3.70 & 2.09 & 0.02 & 1740 & 2091 &  30  & 1.3 & 26.8 & 73  & \\object{HIP 49637}\\\\ & ~9 Mar 2001 & 2451978 & 3.93 & 2.09 & 0.02 &  390 &  777 &  20  & 1.0 & 27.0 & 73  & \\object{31 Leo} \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} Interferometric near-infrared imaging of \\object{IRC+10216} with angular resolutions of better than 100 mas has revealed that its dust shell is clumpy and bipolar, and is changing on a time scale of only $\\sim$1\\,yr (Weigelt et al.\\ \\cite{WeiEtal97}, Weigelt et al.\\ \\cite{WeiEtal98} [hereafter Paper~I], Haniff \\& Buscher \\cite{HanBus98}, Osterbart et al.\\ \\cite{OstEtal00} [hereafter Paper~II], Tuthill et al.\\ \\cite{TutEtal00}). In 1996, four components were identified in the inner dust shell of \\object{IRC\\,+10\\,216} within a radius of 200\\,mas (Weigelt et al. \\cite{WeiEtal97}, Paper~I, Haniff \\& Buscher \\cite{HanBus98}) and were denoted as A, B, C, and D in order of decreasing brightness (see Fig.~\\ref{FKima}). On larger scales the envelope of \\object{IRC\\,+10216} appears to be spherically symmetric (Mauron \\& Huggins \\cite{MauHug99,MauHug00}). Since most dust shells around AGB stars are known to be spherically symmetric on larger scales, whereas most proto-planetary nebulae (PPN) appear in axisymmetric geometry (Olofsson \\cite{Olof96}), it is likely that \\object{IRC\\,+10216} has already entered the transition phase to PPN. This suggests that the break of symmetry already takes place at the end of the AGB evolution. So far, only a few AGB objects are known to show prominent asphericities of their dust shells in the near-infrared, and are therefore believed to have entered this transition phase at the end of their AGB life. This includes, for instance, the carbon star \\object{CIT\\,6} (Monnier et al.\\ \\cite{MonEtal00}), and the oxygen-rich stars \\object{AFGL\\,2290} (Gauger et al.\\ \\cite{GauEtal99}) and \\object{CIT\\,3} (Hofmann et al.\\ \\cite{HofEtal01}).  \\begin{figure*} \\begin{center} \\epsfxsize=44mm \\mbox{\\epsffile{h3683f01.eps}} \\epsfxsize=44mm \\mbox{\\epsffile{h3683f02.eps}} \\epsfxsize=44mm \\mbox{\\epsffile{h3683f03.eps}} \\epsfxsize=44mm \\mbox{\\epsffile{h3683f04.eps}}\\\\[0.3mm] \\epsfxsize=44mm \\mbox{\\epsffile{h3683f05.eps}} \\epsfxsize=44mm \\mbox{\\epsffile{h3683f06.eps}} \\epsfxsize=44mm \\mbox{\\epsffile{h3683f07.eps}} \\epsfxsize=44mm \\mbox{\\epsffile{h3683f08.eps}} \\end{center} \\caption{ $K$-band speckle reconstructions of \\object{IRC\\,+10\\,216} for 8 epochs from 1995 to 2001. The total area is 1\\arcsec$\\times$1\\arcsec. All images are normalized to the brightest pixel and are presented with the same color table. North is up and east is to the left. \\vspace*{0.3cm} } \\label{FKima} \\begin{center} \\epsfxsize=59mm \\mbox{\\epsffile{h3683f09.eps}} \\epsfxsize=59mm \\mbox{\\epsffile{h3683f10.eps}} \\epsfxsize=59mm \\mbox{\\epsffile{h3683f11.eps}} \\\\[0.3mm] \\end{center} \\caption{ $J$-, $H$- and $K$-band speckle reconstructions of \\object{IRC\\,+10\\,216} in March 2001. The total area is 1.6\\arcsec$\\times$1.6\\arcsec. All images are normalized to the brightest pixel and are presented with the same color table. North is up and east is to the left. The tick marks indicate the likely position of the central star. } \\label{FJHKima} \\end{figure*}  Recently, we have developed a 2D radiative transfer model of \\object{IRC\\,+10\\,216} (Men'shchikov et al. \\cite{mbbow2001}; hereafter Paper III) which can explain many aspects of the nebula. In this model, the star is located at component B and is surrounded by an optically thick shell with bipolar cavities. The brightest southern component A is identified with radiation emitted and scattered in the optically thinner southern cavity of the dense shell. However, Weigelt et al. (\\cite{WeiEtal98}), Haniff \\& Buscher (\\cite{HanBus98}), and Tuthill et al. (\\cite{TutEtal00}) argued that the star is located at the position of component A and that components B, C, and D are dust clouds. In the present paper, we attempt to explain the recent changes in the nebula in light of the 2D model of Paper III, and will comment on the alternative interpretations only in passing.  High-resolution near-infrared monitoring of the components A, B, C, and D has already revealed that the dust shell of \\object{IRC\\,+10\\,216} is rapidly evolving. The 6\\,m telescope speckle observations presented in Paper~II cover five epochs between 1995 and 1998 and show that the separations between the different components had steadily been increasing. For example, the distance between the initially brightest components A and B increased by 36\\% during 1995-1998. These results are in very good agreement with Keck telescope $K$-band observations obtained with the highest resolution to date at 7 epochs between 1997 and 1999 by Tuthill et al.\\ (\\cite{TutEtal00}). Such direct observations of the dust-shell evolution offer an ideal opportunity to study the mass-loss process in the late stages of AGB evolution, revealing details of the dust formation process as well as the geometry and clumpiness of the stellar wind.  This paper presents new near-infrared bispectrum speckle interferometry monitoring of \\object{IRC+10216} obtained with the SAO 6\\,m telescope in 1999--2001, adding 3 new epochs to the already available data of dust-shell evolution. The present observations show that the appearance of the dust shell has considerably changed compared to the epochs of 1995 to 1998. This paper is organized as follows. In Sect.~\\ref{Sobs}, $K$-band observations of the dust-shell evolution from 1999 to 2001, a comparison of  $J$-, $H$-, and $K$-band images of  1996/97 and 2001, and $J{-}H$, $J{-}K$, and $H{-}K$ color images are presented. In Sect.~\\ref{Sdis}, these observations are discussed on the basis of the general morphology,                                  % 2D radiative transfer models,            % and dust-formation models.             % Conclusions are given in Sect.~\\ref{Scon}. ",
        "conclusions": "\\label{Sdis} \\subsection{Dust and winds} \\label{Swind} On larger scales, most AGB stars are surrounded by dust shells of spherical symmetry. This holds also for the dust shell of \\object{IRC\\,+10\\,216}  (Mauron \\& Huggins \\cite{MauHug99,MauHug00}). However,  on  sub-arcsecond scales this symmetry appears to break down for giants in very advanced stages of their AGB evolution, with \\object{IRC\\,+10\\,216} being their most prominent representative. The 2D radiative transfer modelling presented in Paper III strongly suggested that the central star of \\object{IRC\\,+10\\,216} is surrounded by an optically thick dust shell with bipolar cavities of a full opening angle of $36^{\\rm o}$ tilted with the southern lobe towards the observer (although there may exist different interpretations, see Sect.~\\ref{Sintro}). The bright compact component A is not the direct light from the central star but is the southern cavity of this bipolar structure dominated by scattered light. According to this model, the carbon star is at the position of the fainter northern component B.  This dust-shell model and the possible existence of acceleration within the cavities may be understood in terms of the beginning operation of interacting winds which later lead to the shaping of planetary nebulae (Kwok \\cite{Kwok82}; Balick \\cite{Bal87}). Within this scenario a fast stellar wind interacts with the fossil, slow AGB wind. The AGB mass-loss is thought to be aspherical and to take preferentially place in the equatorial plane. Accordingly, the polar regions show lower densities than the equatorial regions. Later,   % after a fast wind has developed, the less dense polar material is carved out leading to the formation of outflow cavities. This fast wind can be expected to turn on slowly at the very end of the AGB evolution resulting into the formation of biconical dust-shell structures on sub-arcsecond scales as observed for \\object{IRC\\,+10\\,216}.  The high spatial and temporal resolution of the present observations reveals details of the  mass-loss process even in the immediate vicinity of the dust condensation zone. The monitoring, covering more than  3 pulsation periods, shows that the structural variations are not related to the stellar pulsation cycle in a simple way. This is consistent with the predictions of hydrodynamical models that enhanced dust formation takes place on a timescale of several pulsation periods (Fleischer et al.\\ \\cite{FleiEtal95}). In Sect.~\\ref{Sdust} simplified dust formation models are compared with the  observations. \\subsection{Observed changes interpreted by 2D modeling} \\label{S2dmod}  The large number of epochs available now seems to make the interpretation of the sub-arcsecond structures and their evolution even more difficult and puzzling than before. Here we suggest a physical picture of what presently is going on in this object based on the results of our previous 2D radiative transfer modeling (Paper~III). The model was designed to derive the structure and physical properties of the star and its dusty envelope at a single moment, corresponding to the third epoch (January 23, 1997) of our images. Our simplified attempt to handle the difficult problem of self-consistent {\\em time-dependent} multidimensional radiative transfer modeling will be presented elsewhere (Men'shchikov et al. 2002, submitted).  \\subsubsection{The star and its evolving environment}  Key point to the understanding of the structure and evolution of the sub-arcsecond environment of \\object{IRC\\,+10\\,216} is the knowledge of the position of the central star in the images. As we have suggested in the previous modeling (Paper~III), the brightest compact peak A seen in all images is not the direct light from the central star. The star is most likely located at the position of the fainter peak B, whereas component A is the radiation emitted and scattered in the optically thinner conical cavity of the optically thick bipolar dust shell. The even fainter components C and D in the $H$ and $K$ images were identified with smaller-scale deviations of the density distribution of the circumstellar environment from axial symmetry (Fig.~\\ref{FKima}; Fig.~2 in Paper~III). These inhomogeneities are less opaque than other, more regular regions of the compact dense shell.  An alternative interpretation of the morphology of our images of \\object{IRC\\,+10\\,216} would be that components A and B are the lobes of a bipolar nebula and that the star is located in the dark region {\\em between} them. However, this interpretation disagrees with our radiative transfer modeling presented in Paper~III. A reason for this is that the optically thick dust which would form the bipolar lobes cannot exist so close to the star, at {\\em half} the distance between the components. Thus, the star is probably located either at A or B. The fan-shaped morphology of component A in the $J$- and $H$-band images, as well as other evidence presented in Papers II and III, strongly suggest that the star is located at the position of B.  Having defined the stellar position at component B, we can now interpret the evolving appearance of \\object{IRC\\,+10\\,216} in terms of an increased mass-loss rate during the last $\\sim$ 30 years. One of the main features of the entire 8-epoch sequence of imaging is the relative fading of the stellar component B (Figs.~\\ref{FKima}, \\ref{FJHKima}). Other changes seen in the images are the increasing distances of components A, C, D from the star (Fig.~\\ref{FsepABCD}) and their greatly varying shapes. Our radiative transfer modeling (Paper~III) has shown that the observed components are {\\em cavities} in the dense opaque shell; therefore, the observed motion requires an additional discussion (Sect.~\\ref{Motion}).  \\subsubsection{The near-IR images and model cavities}  In order to better visualize the relative location of the components, we have shown the structure (conical cavities) of our model (Paper~III) in the observed images by dashed lines in Figs.~\\ref{FJHKcont} and \\ref{Fbild}. The lines intersecting at the position of the central star and making an angle of 46{\\degr} between them, outline the biconical geometry of the cavities.  The April 1996 $J$-band image in Fig.~\\ref{FJHKcont} best exhibits the bipolar geometry of \\object{IRC\\,+10\\,216}. Clearly visible at this wavelength are the fan shape of the brightest component A and even the scattered light from the opposite (northern) cavity on the far side of the dense shell. Faint direct stellar light is visible near the origin of the conical cavities. The star (component B) appears much brighter in the January 1997 $H$-band image (Fig.~\\ref{FJHKcont}), whereas the fan shape of the cavity (component A) is less prominent but still well visible at this wavelength. The $H$ and $K$ images in 1997 agree very well with our model presented in Paper~III. The March 2001 $H$ and $K$ images are qualitatively similar to the older images obtained in 1997, but distorted by the appearance of several fainter structures close to component D and by merging of the bright component A with component C. The faint direct stellar light has become more difficult to identify in $H$ and it has been buried in the enhanced dust emission from all the other components in the March 2001 $K$-band image.  The most recent $J$ image of \\object{IRC\\,+10\\,216} with a 3 times higher resolution (50 mas, Fig.~\\ref{FJHKima}) shows in much greater detail the fine structure of the envelope. The image reveals a rather isolated, faint peak at the position of the central star (marked in Fig.~\\ref{FJHKima}). The following evidence suggests that this peak is indeed the direct light from the star, not just a clump that happened to be there: (1) its position angle relative to A (PA $\\approx 20$\\,{\\degr}) is the same as in all the images where the star (component B) is clearly visible, (2) the $H$ image at the same epoch shows a (less isolated) peak at the same position, (3) the distance of approximately 347 mas between the peak and the component A is consistent with the fit of Fig.~\\ref{FsepABCD}.  One might think that the faint stellar peak should be much better visible in $K$ band (Fig.~\\ref{FJHKima}), where optical depths are significantly lower compared to $J$. However, the greatly enhanced hot dust emission in $K$ band may well make it completely invisible. In fact, from the continuum of the model stellar atmosphere used in Paper~III, we estimate that the stellar brightness in $J$ and $K$ bands is approximately the same. The dust model of Paper~III would predict that the star is a factor of $\\sim 1.8 \\times 10^3$ brighter in $K$ than in $J$ band, whereas our new model of \\object{IRC+10216} (Men'shchikov et al., submitted), computed specifically for the latest epoch of March 2001, predicts a larger factor of $\\sim 1.3 \\times 10^4$. On the other hand, from the calibrated color images of Fig.~\\ref{Fbild} we know that the nebula becomes much brighter between $J$ and $K$ (and thus redder), by a factor of $\\sim 1.4 \\times 10^4$. Therefore, it is unlikely that in $K$ the stellar peak is visible better than in $J$ band.  \\subsubsection{Shapes of the near-IR color maps}  The high-resolution $J{-}H$, $J{-}K$, and $H{-}K$ color images in Fig.~\\ref{Fbild} observed in March 2001 confirm our interpretation that the bright components A, C, and D are, respectively, the cavity and smaller-scale inhomogeneities in the dense shell, {\\em not} dense clumps of dust as it might appear from the images alone (Figs.~\\ref{FKima}--\\ref{FJHKcont}). In fact, all the bright components coincide with the blue areas in the color images, the cavity A corresponding to the bluest spot. This is a natural consequence of lower optical depths along those directions from the star, with the ``hot'' stellar photons being scattered into the direction of the observer. This is illustrated in Fig.~\\ref{Fbild} by the dashed lines showing the geometry of our model (Paper~III). The bluest, optically thinnest spot is located precisely inside the conical cavity. The star, obscured by $\\sim$ 40 mag of visual circumstellar extinction, is situated in the red area of the color images, as it was also in the previous epochs (Fig.~16 in Paper~III).  \\subsubsection{Moving dense layer or dust evaporation?} \\label{Motion}  The apparent motion of the components A, C, and D (Fig.~\\ref{FKima}) could be attributed either to the real radial expansion of the opaque dense layer with several ``holes'' in the dense dust formation zone or to a displacement of the dust formation radius due to evaporation of recently formed dust by a hotter environment. Here, we analyze both processes and argue that the temperature-induced displacement of the dust formation zone is acting in \\object{IRC\\,+10\\,216}. For simplicity, we consider a single dust formation radius corresponding to the formation of carbon dust (cf. Paper~III).  One can explain the observed decreasing brightness of the star by assuming a monotonically increasing mass-loss rate and, hence, higher densities and optical depths of the wind in the dust formation zone. Higher mass loss and continuing condensation of new dust in the wind out of the gas phase increase the temperatures of the outflowing gas and dust due to backwarming. Increasing temperatures affect the location of the inner dust boundary of the envelope via dust sublimation, causing its displacement with a velocity that has nothing to do with the outflow motion of the envelope.  Crucial to distinguishing between the real motion of a dense layer and the temperature-induced shift of the dust formation zone are estimates of the apparent velocities of the components A, C, and D relative to the star. For the assumed distance of 130 pc, the linear fit in Fig.~\\ref{FsepABCD} gives for the brightest component A a velocity $v_{\\rm A} \\approx 18$ km\\,s$^{-1}$ in the plane of sky. On the basis of our model (Paper~III), one can derive for the component a deprojected radial velocity $v_{r{\\rm A}} \\approx 19$ km\\,s$^{-1}$.  Since the deprojected radial velocity is higher than the observed (terminal) wind outflow speed in \\object{IRC\\,+10\\,216} of $v \\approx 15$ km\\,s$^{-1}$, it is unlikely that the observed changes reflect just an expansion of a dense layer in which grains are forming. In fact, the standard picture of a stationary stellar wind predicts an acceleration of dust and gas within distances by a factor of $\\sim$ 2--5 larger than the dust formation radius (Steffen et al.\\ \\cite{SteffEtal97}). Due to the radiation pressure on dust grains, the wind velocity increases in this transition zone, approaching asymptotically the terminal velocity at larger distances. As our model (Paper~III) associates the observed components of \\object{IRC\\,+10\\,216} with the dust formation zone, we expect that dust and gas are not yet fully accelerated there, i.e. that the radial outflow velocity $v_r < 15$ km\\,s$^{-1}$.  However, the deprojected velocity $v_{r{\\rm A}} \\approx 19$ km\\,s$^{-1}$ is significantly larger than the expected wind velocity in the dust formation zone. Only if we assume an unlikely lower limit of 100\\,pc (Becklin et al. \\cite{Becklin_etal1969}), does the deprojected velocity $v_{r{\\rm A}}$ approach (from above) $v \\approx 15$ km\\,s$^{-1}$, which is still too high for the dust formation zone. If, however, the actual distance to \\object{IRC\\,+10\\,216} is larger than 130\\,pc, one would obtain $v_r \\ga 19$ km\\,s$^{-1}$ and an even larger discrepancy.  Moreover, there are reasons to believe that the acceleration depicted by the parabolic fit in Fig.~\\ref{FsepABCD} is real and that the actual apparent motion of A is now as fast as $v_{\\rm A} \\approx 26$ km\\,s$^{-1}$. The corresponding deprojected velocity is then $v_{r{\\rm A}} \\approx 28$ km\\,s$^{-1}$, much higher than the expected wind speed, for any realistic distance to \\object{IRC\\,+10\\,216}. Taken in context with increasing optical depths in the shell, this suggests that the observed motions are caused by the rapid dust evaporation due to backwarming and higher temperatures in the dense environment formed by the increased mass loss.  We believe that a reasonable interpretation of the observed changes in \\object{IRC\\,+10\\,216} would be the following picture. During the recent period of increasing mass-loss which started $\\sim$ 20--30 years ago, a compact dense shell has formed around the star. The mass-loss rate was recently as high as $\\dot{M} \\approx 10^{-4} M_\\odot$\\,yr$^{-1}$ (Paper~III) and the innermost regions of the dust shell are expanding outward at a local wind velocity $v \\la 10$ km\\,s$^{-1}$. Dust formation continues in the expanding material, thus increasing its optical depth and obscuring the central star. The optical depths in the polar regions remain significantly smaller than in the other regions of the dense shell, making the cavity (A) and the other components relatively brighter than at previous epochs. As the dense, increasingly optically thick dusty shell expands, steeply rising temperatures inside it (due to the backwarming from the steepening density front) inhibit further dust condensation and evaporate outflowing grains. In effect, these processes have been shifting recently the dust formation radius outward with an average velocity $v_r \\approx 19$ km\\,s$^{-1}$ (or as high as $\\sim 30$ km\\,s$^{-1}$ in 2001, if the apparent acceleration measured in this work is real). One can predict that the star will remain obscured until $\\dot{M}$ starts to drop back to lower values. In a few years from that moment, we could probably be witnessing the star (B) reappearing whereas the cavities becoming relatively fainter. \\subsection{Dust formation models and the fading of B} \\label{Sdust} If the star is located at the position of component B, as is suggested by the two-dimensional radiative transfer model (see Sect.~\\ref{S2dmod} and Paper III), \\nocite{mbbow2001} the fading of B might be caused by the formation of new dust along the line of sight towards the star. To investigate, whether such a scenario would be capable of explaining the observed time scale of the fading of B, we constructed a very simple gas box model. The idea is to follow the process of carbon grain formation in a gas element moving away from the star at a constant velocity. To characterize the changing thermodynamic conditions experienced by the gas element, we assume a power-law gas temperature stratification  \\begin{equation} T^{\\rm g}(r) = T^{\\rm g}_{0} \\left(\\frac{R_0}{r}\\right)^{\\alpha} \\label{tempeq} \\end{equation} and evaluate the gas density structure from mass conservation in a stationary, spherically symmetric configuration \\begin{equation} \\rho^{\\rm g}(r) = \\frac{\\dot{M}}{4 \\pi v}\\frac{1}{r^2} = \\rho^{\\rm g}_{0} \\left(\\frac{R_{0}}{r}\\right)^{2} \\,. \\label{rhoeq} \\end{equation} For the co-moving gas element, the time coordinate is given by \\begin{equation} {\\rm d} t = \\frac{{\\rm d} r}{v}\\,\\,\\,. \\end{equation}  In this gas element, we consider heteromolecular formation and growth of carbon grains which we compute according to the moment method derived by Gail \\& Sedlmayr~(1988)\\nocite{gs88}. The growth process includes reactions with the molecular species C, C$_2$, C$_2$H, and C$_2$H$_2$. In order to calculate the concentrations of the relevant carbon-bearing molecules, chemical equilibrium in the gas phase is assumed. We consider a carbon-rich element mixture with otherwise solar abundances.  \\begin{figure} \\begin{center} \\epsfxsize=8.8cm \\epsffile[18 148 445 664]{h3683f21.eps} \\end{center} \\caption{Temporal evolution of the gas box. The time coordinate is given in units of $P=649$\\,d, the pulsation period of IRC~+10216. {\\it Upper panel}: gas temperature (Eq.~(\\ref{tempeq}), solid line, l.h.s. ordinate) and density (Eq.~(\\ref{rhoeq}), dash-dotted line, r.h.s. ordinate); {\\it Second panel}: dust nucleation rate $J_{*}$ (solid line, l.h.s. ordinate) and number of dust grains per hydrogen nucleus (dash-dotted line, r.h.s. ordinate); {\\it Third panel}: degree of condensation $f_{\\rm cond}$ (solid line, l.h.s. ordinate) and optical depth at 2.2\\,$\\mu$m (dash-dotted line, r.h.s. ordinate); lower panel: normalized emergent intensity at 2.2\\,$\\mu$m (solid line) and observed intensity of component B, normalized to its maximum value (crosses). See text for more details.} \\label{gasboxplot} \\end{figure}  In Fig.~\\ref{gasboxplot} the resulting temporal evolution of the gas box is depicted for the following parameters: $T_{0} = 3000\\,$K at a radius $R_{0} = 3.15 \\times 10^{13}\\,$cm (corresponding to the Stefan-Boltzmann radius for $T_{0}$ and a stellar luminosity of $15000\\,L_{\\odot}$), $\\alpha = 0.5$ (i.e. an optically thin temperature distribution), $\\dot{M} = 10^{-4}M_{\\odot}$yr$^{-1}$ (reflecting a presently increased mass-loss rate of IRC~+10216), a constant outflow velocity of $v = 15\\,$km\\,s$^{-1}$, and a carbon abundance (by number) of $\\epsilon_{\\rm C} = 1.5 \\, \\epsilon_{\\rm O}$.  In the upper panel of Fig.~\\ref{gasboxplot} the prescribed time evolution of gas temperature (Eq.~(\\ref{tempeq})) and density (Eq.~(\\ref{rhoeq})) in the gas box are shown. The second panel depicts the nucleation rate $J_{*}$, i.e. the number of seed particles formed per second and hydrogen nucleus (solid line) and the number of dust grains per hydrogen nucleus (dash-dotted line). In the third panel, the degree of condensation, i.e. the fraction of condensible material that is actually condensed is shown (solid line). The dash-dotted line indicates the optical depth at 2.2\\,$\\mu$m (measured from the star), that would build up in a stationary situation. To calculate the dust extinction coefficient, we assume the small particle limit of Mie theory. This is justified, since the mean particle radius $\\langle a \\rangle$ reached after 15\\,$P$ is $5\\times10^{-2}\\mu$m, giving a maximum size parameter $x = 2 \\pi \\langle a \\rangle / \\lambda = 0.14$ at the considered wavelength of $\\lambda = 2.2\\,\\mu$m. The optical properties of the grains are described by the complex refractive index of amorphous carbon tabulated in Preibisch et~al.~(1993)\\nocite{poyh93}. In the lower panel we plot the normalized intensity at 2.2\\,$\\mu$m that would emerge along the line of sight from the increasingly obscured radiation source (solid line). This intensity is calculated from the formal solution of the radiative transfer equation:  \\begin{equation} \\frac{I}{I_{0}} = e^{-\\tau_{2.2\\mu m}} + \\int_{0}^{\\tau_{2.2\\mu m}}\\frac{B(T(t'))}{I_0}e^{-(\\tau_{2.2\\mu m} - t')} \\,{\\rm d}t' \\,\\,\\,. \\end{equation}  \\noindent Here we assume thermal emission of the dust grains ($B$ is the Kirchhoff-Planck function) and normalize the intensity to $I_{0} = B(T_{0})$.  For comparison, we also plotted in the lower panel of Fig.~\\ref{gasboxplot} the observed intensity of component B, normalized to its intensity in our 1997 speckle image (crosses). The intensity of component B fades from its maximum value to 17\\% of this value within about 1400\\,d, corresponding to 2.2 pulsation periods of IRC~+10216. This time interval is indicated in the lower panel of Fig.~\\ref{gasboxplot} by the vertical dashed lines.  The comparison shows that  dust condensation in front of the star in fact can reproduce the observed time scale of the fading of component B for realistic values of the parameters characterizing our simple toy model.  In Table~\\ref{timetab} the timescale $\\tau$ of the intensity drop to 17\\% of the initial value is summarized for various parameter combinations. We note a steep dependence of $\\tau$ on the C/O ratio, i.e. on the amount of material available for dust formation. Consistent model calculations for dust driven outflows (e.g., Winters et~al.~2000b, Arndt et~al.~1997, Dominik et~al.~1990) \\nocite{wljhs2000,afs96,dgsw90a} show, that the outflow velocity of a carbon-rich dust-driven wind also depends strongly on the C/O ratio. For the observed outflow velocity of \\object{IRC~+10216} of $\\sim 15\\,$km\\,s$^{-1}$, C/O ratios in the range $(1.20 \\la {\\rm C/O} \\la 1.60)$ are required. Therefore, Table~\\ref{timetab} implies, that the present day mass-loss rate of \\object{IRC~+10216} in fact should have increased considerably above the ``canonical'' value of (a few) $10^{-5}M_{\\odot}$/yr derived from CO rotational line observations probing the outer and therefore, older parts of the circumstellar shell (e.g., Sch{\\\"o}ier \\& Olofsson 2001)\\nocite{so2001} or from CO infrared line profiles observed about one decade ago (e.g., Winters et~al.~2000a). \\nocite{wkgs2000}  \\begin{table} \\caption[]{Parameters of the gas box model and resulting fading time scales for an assumed outflow velocity of $v=15\\,$km\\,s$^{-1}$} \\label{timetab} \\begin{flushleft} \\begin{center} \\begin{tabular}{ccccccc} \\hline\\noalign{\\smallskip} $\\alpha$ & $T_{0}$ & $\\rho_{0}$ & $\\dot{M}$    & C/O & $\\tau$ \\\\ &   K   & $10^{-13}\\,$g/cm$^{3}$ & $M_{\\odot}/$yr & & 649\\,d\\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} 0.5      &  2800 & 2.95 & $ 10^{-4}$ & 1.5 & 1.83\\\\%Sascha_max !neue t scale 0.5      &  2800 & 2.95 & $ 10^{-4}$ & 1.4 & $>13$\\\\%Sascha_max 0.5      &  2800 & 2.95 & $ 10^{-4}$ & 1.3 &$\\rightarrow \\infty$\\\\%Sascha_max \\noalign{\\smallskip} 0.5      &  2500 & 4.68 & $ 10^{-4}$ & 1.5 & 0.65\\\\%Sascha_min !neue t scale 0.5      &  2500 & 4.68 & $ 10^{-4}$ & 1.4 & 1.79\\\\%Sascha_min !neue t scale 0.5      &  2500 & 4.68 & $ 10^{-4}$ & 1.3 &$\\rightarrow \\infty$\\\\%Sascha_min \\noalign{\\smallskip} 0.5      &  2500 & 0.47 & $ 10^{-5}$ & 2.1 & 1.58\\\\%Sascha_min !neue t scale 0.5      &  2500 & 0.47 & $ 10^{-5}$ & 2.0 & 3.05\\\\%Sascha_min !neue t scale 0.5      &  2500 & 0.47 & $ 10^{-5}$ & 1.9 &$\\rightarrow \\infty$\\\\%Sascha_min \\noalign{\\smallskip} 0.5      &  3000 & 3.37 & $ 10^{-4}$ & 1.6 & 1.04\\\\%w35 !neue t scale 0.5      &  3000 & 3.37 & $ 10^{-4}$ & 1.5 & 2.22\\\\%present text !neue t scale 0.5      &  3000 & 3.37 & $ 10^{-4}$ & 1.4 & $\\rightarrow \\infty$\\\\ \\noalign{\\smallskip} 0.4      &  3000 & 3.37 & $ 10^{-4}$ & 1.6 & 2.59\\\\ 0.4      &  3000 & 3.37 & $ 10^{-4}$ & 1.5 & 7.53\\\\ \\noalign{\\smallskip} 0.6      &  3000 & 3.37 & $ 10^{-4}$ & 1.5  & 1.15\\\\ 0.6      &  3000 & 3.37 & $ 10^{-4}$ & 1.4  & 5.01\\\\ \\noalign{\\smallskip} 0.7      &  3000 & 3.37 & $ 10^{-4}$ & 1.5 & 0.75\\\\ 0.7      &  3000 & 3.37 & $ 10^{-4}$ & 1.4 & 2.54\\\\ \\noalign{\\smallskip} 0.7      &  3000 & 0.34 & $ 10^{-5}$ & 2.1 & 2.35\\\\ 0.7      &  3000 & 0.34 & $ 10^{-5}$ & 1.9 & 9.62\\\\ 0.7      &  3000 & 0.34 & $ 10^{-5}$ & 1.5 & $\\rightarrow \\infty$\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{flushleft} \\end{table}  Once dust has formed in the circumstellar shell, its opacity would lead to a pronounced backwarming and therefore, to a steeper temperature gradient in the dust formation zone. Accordingly, in our simple gas box model, lower temperatures would be reached at the same density. As a result, dust growth becomes more efficient in this region and a reduced C/O ratio is sufficient to recover the observed fading time scale (see Table~\\ref{timetab}; compare the corresponding entries for $\\alpha = 0.5$ with $\\alpha = 0.6$ and $0.7$).  A more realistic treatment of the dust formation and mass-loss process in \\object{IRC~+10216} would require the simultaneous solution of the time dependent hydrodynamic equations, taking into account the pulsation of the star, together with the equations describing the dust formation process and the radiative transfer problem. Such a consistent investigation will be presented in a forthcoming paper.  An alternative interpretation of the fading of component B, together with its increasing separation from A, could be based on a scenario, where the star is located in the direction of component A, as argued in previous studies by Weigelt et al. (\\cite{WeiEtal98}), Haniff \\& Buscher (\\cite{HanBus98}), and Tuthill et al. (\\cite{TutEtal00}).  In this case, component B would indicate thermal emission from a dust cloud, ejected from the star in a direction almost perpendicular to the line of sight. The temporal evolution of the brightness of B would then be the result of a competition between the initially increasing thermal emission due to ongoing dust formation in the cloud and decreasing temperature of the cloud as it is moving away from the star. Preliminary results for this scenario, which are based on a consistent model calculation, are presented in Winters et~al.~(2002)\\nocite{wbhw2002}. Although the time scale of the fading of component B is reproduced quite well in this model, the calculated intensity ratio between the expanding cloud and the star is smaller than the observed intensity ratio between components B and A by about a factor of 3."
    },
    "0206/astro-ph0206145_arXiv.txt": {
        "abstract": "Recent wide-field photometric surveys (Lee et al.; Pancino et al.) have shown the existence of a previously unknown metal-rich ($[Fe/H]\\simeq -0.6$) stellar population in the galactic globular cluster $\\omega$ Centauri.  The discovery of this new component, which comprises only a small percentage ($\\sim 5\\%$) of the entire cluster population, has added a new piece to the already puzzling picture of the star formation and chemical evolution of this stellar system.  In this Letter we show that stars belonging to the newly discovered metal-rich population have a coherent bulk motion with respect to the other cluster stars, thus demonstrating that they formed in an independent self-gravitating stellar system.  This is the first clear-cut evidence that extreme metal-rich stars were part of a small stellar system (a satellite of $\\omega$ Centauri?) that has been accreted by the main body of the cluster. In this case, we are witnessing an {\\em in vivo} example of hierarchical merging on the sub-galactic scale. ",
        "introduction": "\\label{intro}  The giant globular cluster $\\omega$~Centauri is the only Galactic globular that shows multiple stellar populations, with a very different heavy metal content. Large photometric \\citep{lee,hk00,hughes,p00} and spectroscopic \\citep{nor96,smith} surveys have demonstrated the presence of a main population at ${\\rm[Fe/H]}=-1.6$ with a metal-rich tail extending up to ${\\rm[Fe/H]}=-0.6$. The detailed abundance patterns of tens of red giants spanning the entire metallicity range have been derived \\citep[see][]{norris,smith,p02}. The cluster also shows dynamical anomalies with respect to ordinary globulars, since it is partially substained by rotation \\citep{may97,merritt}. Moreover, while the metal-poor giants (${\\rm[Ca/H]}<-1.2$) do rotate, the metal-rich ones do not show any sign of rotation \\citep[][see below for further discussion]{nor97}.  Since its vicinity allows a very detailed study of its stellar content, this giant cluster represents a cornerstone in our understanding of the formation, the chemical enrichment and the dynamical evolution of  stellar systems. However, in spite of the huge observational effort carried out so far, the global scenario for the formation and evolution of $\\omega$~Cen still remains a mystery \\cite[see for instance][and discussions therein]{nor96,smith}. To explain its peculiar properties and its heterogeneous mix of populations, different hypotheses have been suggested: {\\it (i)} $\\omega$~Cen is the relic of a larger galaxy partially disrupted by the tidal field of the Milky Way \\citep{freeman,lee,maj}; {\\it (ii)} it is the result of the merging of two globulars \\citep{searle,icke,makino}.  However, while some hints of a complex dynamical history have been found \\citep{nor97,jur}, none of the evidence presented so far is conclusive, and most important, the simple model of the merging of two ordinary globulars cannot fit the observed broad metallicity distribution \\citep{nor96,smith}.    \\subsection{The sub-populations of $\\omega$ Centauri} \\label{sub-pops}   Recent wide-field photometric studies have added a new piece to this already complex jigsaw puzzle: an additional, well-defined Red Giant Branch (RGB) sequence has been discovered to the red side of the main RGB in the Color Magnitude Diagram \\cite[CMD;][]{lee,p00}.  The discovery of this anomalous branch (hereafter RGB-a) unveiled the existence of a further, previously unknown, metal rich sub-population in $\\omega$~Centauri.  In particular \\citet{p00}, taking advantage of the high accuracy and sensitivity of their $(B,B-I)-$CMD, have identified three main sub-populations on the basis of their photometric properties: the main metal-poor component (MP; ${\\rm[Fe/H]}<-1.4$); the metal-intermediate population (MInt; $-1.4\\le{\\rm[Fe/H]}<-0.8$) and the anomalous population (RGB-a) which represents the extreme metal-rich end of the observed metallicity distribution.   The approximate CMD locations of the three sub-populations defined above are shown in Figure~\\ref{wfi-cmd}.  Stars belonging to the RGB-a are plotted as {\\it small solid triangles}.  Note that only a few stars belonging to this population were known before the discovery of the anomalous branch: in particular, only 6 RGB-a stars (marked as {\\it large open triangles} in Figure~\\ref{wfi-cmd}) have previous metal abundance determinations \\citep{evans,nor96} and only 8 have radial velocity determinations \\citep{may97}.  Though the RGB-a population comprises only a small fraction \\citep[$\\sim5\\%$;][]{p00} of the cluster's stellar content, its origin and evolution is the new challenging question in the picture of $\\omega$~Cen formation and evolution.  As part of a long-term project \\citep{post} devoted to reconstruct the global evolutionary history of $\\omega$~Centauri, we are building a comprehensive catalog combining all the relevant information available in the literature for stars in this cluster. Here we present the first results obtained by cross-correlating our wide field, multi-band photometric catalog \\citep{p00}, comprising more than 220,000 stars, with the recently published proper motion study for $\\sim 10,000$ stars brighter than $V\\sim16.0$ \\citep{vanl00}. The synthesis of these data-sets allows us to study for the first time the kinematical properties of the RGB-a stellar population. ",
        "conclusions": "By cross-correlating the photometric catalog by \\citet{p00} and the proper motion catalog by \\citet{vanl00}, we have shown that {\\it (i)} the proper motion distribution of the RGB-a population is not compatible with the one of the dominant (MP) population; {\\it (ii)} the proper motion distribution of the RGB-a is significantly different also from the MInt population \\citep[which includes the metal-rich group by][]{nor97}, {\\it (iii)} the radial velocity ensures us that the coherent moving RGB-a group is gravitationally tied to the main body of $\\omega$~Cen.  This strongly suggest that the RGB-a sub-population had an (at least partially) independent origin with respect to the bulk of the cluster population (i.e., MP and/or Mint stars).  This conclusion is also supported by other observational results.  For example, \\citet{p00} showed that the spatial distribution of the RGB-a stars is elongated perpendicularly to the major axis of the dominant MP population of $\\omega$~Centauri. Furthermore, both high-resolution optical \\citep{p02} and medium-resolution IR \\citep{o02} spectroscopic analyses have shown that RGB-a stars are less $\\alpha$-enhanced than the other sub-populations in $\\omega$~Cen. This fact suggests that the interstellar medium from which they formed could have been polluted by SNIa ejecta, as opposed to {\\em all} the other cluster stars \\cite[see, e.g.][and references therein]{cunha}. In addition, it is worth noticing that all of the six RGB-a stars studied before (large empty triangles in Figure 1) exhibit strong BaII lines  \\citep{evans} and in gerneral very high abundances of post-iron peak elements \\citep{evans,vws02}. These elements are thought to be produced mainly by intermediate-low mass asymptotic giant brach stars \\citep{busso}.  On the other hand, many authors \\citep[see e.g.][and references therein]{norris,smith} showed that the s-process elements overabundance seems to increase continuously with [Fe/H], at least for stars belonging to the RGB-MP and MInt populations, since only few data are available for the RGB-a population. For example, according the abundance analysis by \\citet{vws02}, ROA300 (an RGB-a member) shows a substantially higher abundance of s-process elements with respect to star ROA201 (which is conversely an MInt member). If this kind of discontinuous behaviour is confirmed, it will further support the hypothesis that the RGB-a population had a very different evolutionary and chemical history with respect to the other stars of $\\omega$~Centauri.      In the simplest conceivable scenario, the RGB-a stars originated as an independent stellar system, now trapped and disrupting in the potential well of $\\omega$~Centauri. Since the chance of capturing a completely independent object orbiting in the halo of the Galaxy is highly unlikely, it is reasonable to presume that the accreted system was a former satellite of $\\omega$~Centauri.  This fact would be in agreement with the hypothesis that $\\omega$~Centauri is in fact the relic of a larger galaxy, partially disrupted by the Galactic tidal field \\cite[as suggested by][]{freeman,lee,dinescu,maj}.  If this is the case, we are witnessing the process of hierarchical merging, simultaneously occurring on two very different scales: the $\\omega$~Centauri system is merging into the Milky Way and the RGB-a system is merging into $\\omega$~Centauri.  High resolution Cold Dark Matter (CDM) models \\citep{ben} predict that hierarchical clustering processes do occur down to the dwarf-galaxy scale ($M \\sim 10^7 ~M_{\\odot}$) even at the present day.  Apparently, we have found an {\\em in vivo} example, suggesting that hierarchical merging occurs also on sub-galactic scales ($M \\sim 10^{4-5} ~M_{\\odot}$).  Without a detailed knowledge of the structural parameters for the RGB-a system, no quantitative conclusion can be made concerning the timescales involved in the accretion of the RGB-a by $\\omega$~Cen. However, as an approximate criterium, we can assume that an accreted self-gravitating stellar system can resist the tidal field stress of the main accreting system if its density is larger than the average density of the host system at its orbital radius \\citep{oh,andy}. From the data at our disposal \\citep{p00,post,russo}, we can roughly estimate that the RGB-a population has a central density comparable to that of the MP population and a {\\it core radius} significantly smaller ($r_c^{RGB-a}\\sim 0.5 r_c^{MP}$) than that of the main MP population. Then, according to equation~18 by \\citet{gerh}, the RGB-a system is expected to dissolve under the tidal strain of $\\omega$~Cen if its orbital radius is lower than $\\sim1r_c^{MP}$, i.e. the ``dangerous'' zone for the accreted RGB-a should be limited to the {\\it core} of $\\omega$~Cen.  The observational facts collected so far \\citep{nor97,jur,lee,p00,post} point towards a multiple merging and/or accretion event in the past history of $\\omega$~Cen. However unlikely this may appear, we have no other way to explain the structural and kinematical properties of the various cluster sub-populations. We may note that, as a typical galaxy evolves by accreting smaller systems (i.e., the Sagittarius Dwarf Spheroidal for the Milky Way), the past history of this giant cluster (or dwarf galaxy?) was probably characterized by multiple accreting/merging episodies. Some signatures of these remote events are still observable in the relic of the distrupting sub-systems. The exact connection between the different accreted stellar components (MInt and RGB-a), still remains to be fully understood. A complete radial velocity study of the more metal rich components in $\\omega$~Cen will shed more light on this point."
    },
    "0206/astro-ph0206373_arXiv.txt": {
        "abstract": "We present Westerbork Synthesis Radio Telescope H\\,{\\sc i} images, Lovell telescope multibeam H\\,{\\sc i} wide-field mapping, William Herschel Telescope longslit echelle Ca\\,{\\sc ii} observations, Wisconsin H$\\alpha$ Mapper (WHAM) facility images, and IRAS ISSA 60 and 100 micron coadded images towards the intermediate velocity cloud (IVC) at +70 km\\,s$^{-1}$, located in the general direction of the M\\,15 globular cluster. When combined with previously-published Arecibo data, the H\\,{\\sc i} gas in the IVC is found to be clumpy, with a peak H\\,{\\sc i} column density of $\\sim$ 1.5 $\\times$ 10$^{20}$ cm$^{-2}$, inferred volume density (assuming spherical symmetry) of $\\sim$ 24 cm$^{-3}$ / $D$ (kpc), and maximum brightness temperature at a resolution of $81^{\\prime\\prime} \\times 14^{\\prime\\prime}$ of $\\sim$ 14 K. The major axis of this part of the IVC lies approximately parallel with the Galactic plane, as does the low velocity H\\,{\\sc i} gas and IRAS emission. The H\\,{\\sc i} gas in the cloud is warm, with a minimum value of the full width half maximum (FWHM) velocity width of 5 km\\,s$^{-1}$ corresponding to a kinetic temperature, in the absence of turbulence, of $\\sim$ 540 K. From the H\\,{\\sc i} data, there are indications of two-component velocity structure. Similarly, the Ca\\,{\\sc ii} spectra, of resolution 7 km\\,s$^{-1}$, also show {\\em tentative} evidence of velocity structure, perhaps indicative of cloudlets. Assuming there are no unresolved narrow-velocity components, the mean values of log$_{10}$(N(Ca\\,{\\sc ii} K) cm$^{-2}$) $\\sim$ 12.0 and Ca\\,{\\sc ii}/H\\,{\\sc i} $\\sim$ 2.5$\\times$10$^{-8}$ are typical of observations of high Galactic latitude clouds. This compares with a value of Ca\\,{\\sc ii}/H\\,{\\sc i} $>$ 10$^{-6}$ for IVC absorption towards HD 203664, a halo star of distance 3 kpc, some 3.1 degrees from the main M\\,15 IVC condensation. The main IVC condensation is detected by WHAM in H$\\alpha$ with central LSR velocities of $\\sim$ 60$-$70 km\\,s$^{-1}$, and intensities uncorrected for Galactic extinction of up to 1.3 Rayleigh, indicating that the gas is partially ionised. The FWHM values of the H$\\alpha$ IVC component, at a resolution of 1 degree, exceed 30 km\\,s$^{-1}$. This is some 10 km\\,s$^{-1}$ larger than the corresponding H\\,{\\sc i} value at similar resolution, and indicates that the two components may not be mixed. However, the spatial and velocity coincidence of the H$\\alpha$ and H\\,{\\sc i} peaks in emission towards the main IVC component is qualitatively good. If the H$\\alpha$ emission is caused solely by photoionisation, the Lyman continuum flux towards the main IVC condensation is $\\sim$ 2.7$\\times$10$^{6}$ photons cm$^{-2}$ s$^{-1}$. There is not a corresponding IVC H$\\alpha$ detection towards the halo star HD 203664 at velocities exceeding $\\sim$ 60 km\\,s$^{-1}$. Finally, both the 60 and 100 micron IRAS images show spatial coincidence, over a 0.675$^{\\circ} \\times$0.625$^{\\circ}$ field, with both low and intermediate velocity H\\,{\\sc i} gas (previously observed with the Arecibo telescope), indicating that the IVC may contain dust. Both the H$\\alpha$ and tentative IRAS detections discriminate this IVC from High Velocity Clouds although the H\\,{\\sc i} properties do not. When combined with the H\\,{\\sc i} and optical results, these data point to a Galactic origin for at least parts of this IVC. ",
        "introduction": "The study of intermediate velocity clouds (IVCs) remains one of the most challenging in contemporary Galactic astronomy, with several issues concerning IVCs remaining unresolved. These include, but are not limited to, the method of their formation, their relationship (if any) with high velocity clouds (HVCs), and the question as to whether IVCs are sites of star formation in the halo of the Galaxy (Kuntz \\& Danly 1996; Christodoulou, Tohline \\& Keenan 1997; Ivezic \\& Christodoulou 1997). This latter possibility is underpinned by the fact that within the Galactic halo, there exists a population of early B-type stars whose velocities, ages and distances from the Galactic plane ($z$) are incompatible with them being formed within the disc. A possible site for their formation is IVCs/HVCs via cloud-cloud collisions and subsequent compression of the gas (Dyson \\& Hartquist 1983). Such collisions are thought to be a viable star formation mechanism within at least the discs of galaxies, albeit where the gas density and cloud-cloud collision rates are somewhat higher than inferred in IVCs/HVCs (Tan 2000).  The solution to both the star formation question and also any possible relationship between HVCs and IVCs requires both the analysis of aggregate parameters of well-defined samples of IVCs and HVCs, and also more detailed studies of individual objects. In this paper we report on radio H\\,{\\sc i} aperture synthesis, H\\,{\\sc i} multibeam wide field mapping, longslit Ca\\,{\\sc ii} observations, Wisconsin H$\\alpha$ Mapper (WHAM) facility images, and IRAS sky-survey archive data retrieval towards a particular IVC located in the general direction of the M\\,15 globular cluster (RA=21$^{\\rm h}$ 29$^{\\rm m}$ 58.29$^{\\rm s}$, Dec=+12$^{\\circ}$ 10$^{\\prime}$ 00.5$^{\\prime\\prime}$ (J2000); $l$=65.01$^{\\circ}$, $b$=--27.31$^{\\circ}$). These observations are amongst the first H\\,{\\sc i} synthesis data to be taken of positive-velocity IVCs, which remain poorly-studied as a group of objects.  The M\\,15 H\\,{\\sc i} cloud lies at a velocity of $\\sim$ +70 km\\,s$^{-1}$ in the dynamical Local Standard of Rest (Allen 1973); its distance tentatively lies between $\\sim$ 0.8--3 kpc (Little et al. 1994; Smoker et al. 2001a). The upper distance limit is gleaned from the fact that IVC absorption at $\\sim$+70 to +80 km\\,s$^{-1}$ is observed in the spectrum of HD 203664, a halo star of distance 3 kpc and $\\sim$ 3.1 degrees from M\\,15 (Little et al. 1994), combined with the detection of IV H\\,{\\sc i} approximately mid-way between the M\\,15 IVC and HD 203664.  The deviation velocity of the M\\,15 IVC at such a mid-Galactic latitude puts it in on the borderline between the normal definitions for intermediate and high velocity clouds (c.f. Fig. 1 of Wakker 1991), although in common with Sembach (1995) and Kennedy et al. (1998) here we classify it as an IVC. The line-of-sight position of the M\\,15 IVC is between the negative-velocity Local Group barycentre cloud Complex G and the Galactic centre clouds (Fig. 8 of Blitz et al. 1999), hence the M\\,15 IVC is a part of IVC Complex gp (Wakker 2001).  Previous observations in H\\,{\\sc i} emission using the Lovell and Arecibo telescopes (Kennedy et al. 1998; Smoker et al. 2001a) have shown that the IVC consists of several condensations of gas spread out over an area of more than 3 square degrees, with structure existing down to the previous resolution limit of $\\sim$ 3 arcmin. The brightest component is located towards M\\,15 itself and has peak H\\,{\\sc i} column density at the Arecibo resolution of $\\sim$ 8$\\times$10$^{19}$ cm$^{-2}$. In this paper, we study this part of the IVC at higher spatial resolution. The mass of this particular clump is $\\sim$ 20 $D^{2}$ M$_{\\odot}$, (where $D$ is the distance in kpc), thus for this particular object, in the absence of H$_{2}$, there is insufficient neutral gas to form an early-type star. Low-resolution absorption-line Ca\\,{\\sc ii} and Na\\,{\\sc i} spectroscopy (Lehner et al. 1999) towards cluster stars tentatively found cloud structure (or variations in the relative abundance) over scales as small as a few arcsec, with fibre-optic array mapping in the Na\\,{\\sc i} D absorption lines (Meyer \\& Lauroesch 1999) obtaining similar results with structure visible on scales of $\\sim$ 4 arcsec (velocity resolution $\\sim$ 16 km\\,s$^{-1}$). Using empirical relationships between the sodium and hydrogen column densities, Meyer \\& Lauroesch (1999) derived values of the H\\,{\\sc i} column density towards the cluster centre of $\\approx$ 5$\\times$10$^{20}$ cm$^{-2}$, some 6 times higher than that found using the Arecibo H\\,{\\sc i} data alone; the difference may be attributable to fine-scale cloud structure. Assuming spherical symmetry, a volume density of $\\approx$ 1000 cm$^{-3}$ is implied by these latter results, similar to values obtained for gas in the local ISM (e.g. Faison et al. 1998, although see Lauroesch, Meyer \\& Blades 2000). Such a high volume density and implied overpressure with respect to the ISM perhaps indicates that the assumption of spherical symmetry is invalid and that there may be some sheet like geometry in the IVC as has been postulated for low-velocity gas (Heiles 1997).  In the current paper, we extend our studies of the M\\,15 IVC to higher resolution and different wavelength regions in order to investigate three areas. Firstly, H\\,{\\sc i} synthesis mapping, WHAM H$\\alpha$ and IRAS ISSA survey data retrieval towards the IVC were performed to see if the H\\,{\\sc i}, H$\\alpha$ and infrared properties of the M\\,15 IVC are compatible with either low velocity gas or HVCs in general, and whether there are any differences between the types of object, perhaps attributable to differences in the formation mechanism or current environment (for example, distance from the ionising field of the Galaxy). Secondly, wide-field medium-resolution H\\,{\\sc i} data were obtained in search of more IVC components, to trace how the gas kinetic temperature changes with sky position, and possibly determine the relative distance of cloud components from the Galactic plane (c.f. Lehner 2000). Thirdly, longslit echelle Ca\\,{\\sc ii} observations were undertaken, using the centre of M\\,15 as a background continuum source, in order to look for small-scale velocity and column density substructure within the IVC which could indicate the presence of cloudlets, collisions between which in certain IVCs may be responsible for star formation in the Galactic halo.  Section \\ref{observations} describes the observations and data reduction, Sect. \\ref{results} gives the results, Sect. \\ref{disc} contains the discussion and Sect. \\ref{concl} presents a summary and the conclusions. ",
        "conclusions": "\\label{concl}  The current H\\,{\\sc i} WSRT synthesis observations have shown that on scales down to $\\sim$ 1 arcmin, the M\\,15 IVC shows substructure, with variations in the column density of a factor of $\\sim$ 4 on scales of $\\sim$ 5 arcmin being observed, corresponding to scales of $\\sim$ 1.5 $D$ pc, where $D$ is the the IVC distance in kpc. Of course, this is not an unexpected finding, but once again demonstrates that great care must be taken in interpreting quantities such as cloud metallicities which are derived from a combination of low-resolution radio plus optical data. The Lovell telescope H\\,{\\sc i} observations towards this cloud demonstrated how relatively large areas of sky can be mapped with the multibeam system in a short period of time in the search for IVCs and HVCs. These data showed that the M15 IVC has components spread out over several square degrees, with component `D' being mapped for the first time at medium resolution (12 arcmin) and having a similar column density to the IV gas centred upon M15 itself. Both the H\\,{\\sc i} emission-line and Ca\\,{\\sc ii} absorption-line data showed {\\em tentative} evidence for velocity substructure, perhaps indicative of cloudlets. The Ca\\,{\\sc ii}/H\\,{\\sc i} value of $\\sim$ 2.5$\\times$10$^{-8}$ towards the main M\\,15 condensation is similar to that previously observed in other IVCs and HVCs. Towards HD 203664, the {\\em observed} lower limit of 10$^{-6}$ is somewhat higher, although this may be caused by factors such as the H\\,{\\sc i} beam being unfilled or partial ionisation of the gas on this sightline. The H\\,{\\sc i} properties of the M\\,15 IVC are indistinguishable from HVCs, although with the lack of distance information towards most HVCs, comparisons are difficult.  The tentative detection of infrared emission from the M15 IVC, as in other IVCs, does distinguish it from HVCs, and either points to the M15 IVC containing more dust, and/or being closer to the heating field of the Galactic plane than HVCs, which as a class of objects are not detected in the IRAS wavebands. Similarly, the relatively strong H$\\alpha$ emission (exceeding 1 Rayleigh) towards parts of the M15 IVC, {\\em if} caused by photoionization, may place it closer to the Galaxy than HVCs. Again, however, this finding is uncertain due to the problem in distinguishing photoionisation from shock ionisation, uncertainties in dust content, and differences in H\\,{\\sc i} volume densities in different objects studied thus far.  Future work towards this cloud should include higher-signal-to-noise observations in the Ca\\,{\\sc ii} line in order to determine if the cloud velocity substructure tentatively found in the current observations is in fact real, and whether the Ca\\,{\\sc ii}/H\\,{\\sc i} ratio determined by the current observations is lower than towards the HD 203664 sightline. This should be combined with $^{12}$CO(1--0) sub-mm observations in order to determine if molecular material exists towards the peaks in H\\,{\\sc i} column density and out of which stars may form. The determination of the falloff in H\\,{\\sc i} column density, of the cloud to low column density limits would also indicate the ionisation properties of the object and whether or not there is any interaction between the M\\,15 IVC gas and low velocity material. Finally, UV observations towards M\\,15 globular cluster stars, although difficult, would provide important information on the absolute metallicity of the gas towards this object for comparison with the HD 203664 sightline."
    },
    "0206/astro-ph0206003_arXiv.txt": {
        "abstract": "I review the arguments motivating models for massive star formation via stellar collisions. I then describe how the standard accretion scenario, involving the collapse of a quasi-hydrostatic gas core, can produce high-mass stars in the pressurized regions of forming star clusters. I argue that the observational evidence, particularly in the Orion hot core, favors the standard accretion paradigm. ",
        "introduction": "Two basic models for how a massive star accumulates its mass are debated. The conventional case, which I refer to as the {\\it standard accretion model}, involves the inside-out collapse of a gravitationally bound, pre-stellar gas core from approximate hydrostatic equilibrium. A small amount of angular momentum creates an accretion disk, via which mass reaches the star in an ordered manner. Low-mass stars appear to form in this way (Shu, Adams, \\& Lizano 1987).  The {\\it collisional model} involves the coalescence of smaller stars or protostars in dense stellar clusters (Bonnell, Bate, \\& Zinnecker 1998). These lower-mass stars may first form via standard accretion from small gas cores that have fragmented from the protocluster gas clump. They may also build up their mass via Bondi-Hoyle accretion from the clump. ",
        "conclusions": "The collisional model for star formation can account qualitatively for some of the observational properties of massive stars, such as their tendency to form in the centers of clusters. However, it is difficult to achieve the necessary stellar densities ($\\sim 10^8\\:{\\rm pc^{-3}}$) for this process to be efficient. The standard accretion model, modified to account for the high pressures and turbulent, nonthermal conditions of massive star-forming clumps can achieve the high-accretion rates necessary to overcome radiation pressure and achieve short formation timescales. Crowding is not a serious problem. Nearby massive star-forming regions show signatures of disks and collimated outflows, suggesting the accretion picture is relevant to the formation of stars with masses up to at least $\\sim 20-30\\sm$."
    },
    "0206/astro-ph0206235_arXiv.txt": {
        "abstract": "In a recent Letter, Bergh\\\"ofer \\& Bowyer rediscussed the analysis of BeppoSAX LECS data of the cluster of galaxies Abell~2199 as presented by Kaastra et al., in particular the detection of a soft X-ray excess.  Bergh\\\"ofer \\& Bowyer stated that their analysis method is better suited and does not show evidence for a soft X-ray excess.  Here we find it necessary to publish a rebuttal, because it can be demonstrated that the method used by Bergh\\\"ofer \\& Bowyer is oversimplified, leading to an erroneous result.  As a consequence, their statement that our initial analysis is incorrect is invalid and the detection of a soft X-ray excess in A~2199 is still confirmed. ",
        "introduction": "The X-ray spectrum of the cluster of galaxies Abell\\,2199 has been studied by many instruments.  The detection of a soft X-ray excess in this cluster was first claimed by \\citet{bowyer98}, based upon {\\it Extreme Ultraviolet Explorer (EUVE)} data, although they do not represent any analysis work. \\citet{kaastra99} analyzed the {\\it BeppoSAX} data of this cluster of galaxies and found evidence of both a soft and a hard X-ray excess at radii larger than 300\\,kpc.  This last analysis was based on spatially resolved spectroscopy with data from the {\\it BeppoSAX}, {\\it EUVE} and {\\it ROSAT} missions.  However, in a recent Letter, \\citet{berghoefer02}, (hereafter BB) made categorical statements that the analysis of Kaastra et al.  is flawed.  BB say explicitly that \"Unfortunately, the telescope sensitivity profile used is likely to be incorrect,\" that \"{\\it BeppoSAX} LECS does not detect an EUV excess when the data are analyzed correctly,\" that \"using a procedure better suited to the analysis of extended sources we show that there is no excess in Abell~2199,\" and that \"these findings appeared to support the (incorrect) finding of an excess in this cluster using {\\it EUVE} data.\"  In this Letter, we discuss the recent conclusions of BB, and we show that BB in fact did not use a different \"telescope sensitivity profile\" and that their conclusions are not based upon improved calibration knowledge.  Therefore, in this Letter, we demonstrate that by pursuing their own highly nonstandard approach to {\\it BeppoSAX} data, the results of BB are susceptible to many problems.  For a clear discussion of this controversy, we focus on the data analysis procedures and omit the discussion on the physical implications of the existence of a soft excess, since this has already been presented in our earlier work \\citep{kaastra99}. ",
        "conclusions": "We have shown in this Letter that the analysis method of BB for assessing the soft excess in A~2199 is inadequate and leads to erroneous conclusions regarding the presence of a soft excess.  In fact, the method of BB is oversimplified since it neglects the intricacies associated with the position and energy dependence of the effective area as well as the spatial/spectral dependence of the LECS PSF.  As a consequence of their inadequate analysis, the results presented by BB are misleading, in the sense that a simple method which makes use of radial profile ratios is purported to be better than a full, sophisticated spatial/spectral analysis, without assessing both methods in a controlled experiment.  As we have shown here their method leads to unpredictable, erroneous results.  The discussion of BB is also misleading because they claim that the finding of a soft excess by Kaastra et al.  is due to calibration problems, while BB neither established this nor used a better LECS calibration.  Finally we have shown that the method of BB is also wrong because of a misunderstanding of the {\\it BeppoSAX} background and corresponding background subtraction errors.  In conclusion, we have shown in this Letter that our previous results on the soft X-ray excess in A~2199 are still valid within the uncertainties related to the spectral and spatial sensitivity of the {\\it BeppoSAX} LECS detector.  It is clear that a further investigation of the presence of the soft excess in A\\,2199, confirming or refuting it, can only come either from a thorough recalibration of the BeppoSAX instruments or from forthcoming observations with other satellites with higher sensitivity, such as the {\\it XMM-Newton} satellite."
    },
    "0206/nucl-ex0206014_arXiv.txt": {
        "abstract": "The \\pa\\ reaction was recognized as one of the most important for gamma ray astronomy in novae as it governs the early 511~keV emission. However, its rate remains largely uncertain at nova temperatures. A direct measurement of the cross section over the full range of nova energies is impossible because of its vanishing value at low energy and of the short \\fo\\ lifetime. Therefore, in order to better constrain this reaction rate, we have performed an indirect experiment taking advantage of the availability of a high purity and intense radioactive \\fo\\ beam at the Louvain~La~Neuve RIB facility. We present here the first results of the data analysis and discuss the consequences. ",
        "introduction": "Gamma--ray emission from classical novae is dominated, during the first hours, by positron annihilation resulting from the beta decay of radioactive nuclei. The main contribution comes from the decay of \\fo\\ (half--life of 110~mn) and hence is directly related to \\fo\\ formation during the outburst. (See the astrophysical discussions in references \\cite{Gom98,Her99,F00} and by Hernanz in these proceedings.) A good knowledge of the nuclear reaction rates of production and destruction of \\fo\\ is required to calculate the amount of \\fo\\ synthesized in novae and the resulting gamma--ray emission. The rate (see ref.~\\cite{WK82}) relevant for the main mode of \\fo\\ destruction (i.e, through \\pa) has been the object of many recent experiments\\cite{Gra00,Bar01} (see also Bardayan in these proceedings and refs. in \\cite{F00}). However, this rate remains poorly known at nova temperatures (lower than 3.5\\power{8}~K) due to the scarcity of spectroscopic information for levels near the proton threshold in the compound nucleus \\nen. This uncertainty is directly related to the unknown proton widths ($\\Gamma_p$) of the first three levels ($E_x$, $J^\\pi$ = 6.419~MeV, 3/2$^+$; 6.437~MeV, 1/2$^-$ and 6.449~MeV, 3/2$^+$). The tails of the corresponding resonances (at respectively $E_R$ = 8~keV, 26~keV and 38~keV) can dominate the astrophysical factor in the relevant energy range\\cite{F00}. As a consequence of these nuclear uncertainties, the \\fo\\ production in nova and the early gamma--ray emission is uncertain by a factor of 300\\cite{F00}. This supports the need of new experimental studies to improve the reliability of the predicted annihilation gamma--ray fluxes from novae. ",
        "conclusions": ""
    },
    "0206/astro-ph0206273.txt": {
        "abstract": "{ The effects of gas accretion on spiral disk dynamics and stability are studied through N-body simulations, including star formation and gas/stars mass exchange. The detailed processes of bar formation, bar destruction and bar re-formation are followed, while in the same time the disk to bulge ratio is varying. The accreted gas might be first prevented to flow inwards to the center by the bar gravity torques, which maintains it to the outer Lindblad resonance. While the first bar is weakening, the accreted gas replenishes the disk, increasing the disk-to-bulge ratio, and the disk self-gravity. A second bar is then unstable, with a higher pattern speed, due both to the increased mass, and shorter bar length. Three or four bar episodes have been followed over a Hubble time. Their strength is decreasing with time, while their pattern speed is increasing. Detailed balance of the angular momentum transfer and evolution can account for these processes. The gas recycled through star formation, and rejected through stellar mass loss plays also a role in the disk dynamics. Implications on the spiral galaxy dynamics and evolution along the Hubble sequence, and as a function of redshift are discussed. ",
        "introduction": "Bars are an essential feature in galaxy evolution. About two-thirds of spiral galaxies are barred (de Vaucouleurs 1963), one third being strongly barred (SB), the other third mildly barred (SAB). In the near-infrared, where the central structure is unveiled from dust, the fraction is even higher (Eskridge et al. 2000).  In the last decade, it has been realized that bars are not long lived features, in particular in gaseous spiral disks, and that galaxies are not frozen for a Hubble time in their morphological class: a bar can be destroyed by large radial gas inflow and mass accumulation in the center (Hasan \\& Norman 1990, Pfenniger \\& Norman 1990, Hasan et al. 1993, Friedli \\& Benz 1993). The mass concentration destroys the orbital structure that supported the bar, and this begins through the creation of two strong inner Lindblad resonances (ILR). The process is initiated by the strong gravity torques exerted by the bar on the gaseous spiral arms, which decrease afterwards as the bar weakens. It can lead to the decoupling of a secondary bar, embedded into the primary, just inside its ILR (Friedli \\& Martinet 1993, Combes 1994), and this is likely to weaken the primary bar.  The bar phenomenon can then be a self-regulated process (e.g. Combes 2000). A tantalizing scenario is that a galaxy may have several bar episodes in its life, and that the fraction of barred galaxies in the de Vaucouleurs classification only reflects the time percentage that a spiral galaxy spends as a barred object. How can a bar be revived after a destruction or weakening event? The disk self-gravity must be enhanced to counteract the stabilizing influence of the bulge and central mass concentration, enforced by the first bar episode. This may be obtained through external gas accretion, that settles in the disk. Significant amounts of matter can be accreted in the life-time of the galaxy (e.g. Katz et al. 1996), and we consider realistic that the mass of a galaxy could double in a few Gyrs. The presence of gaseous warps in practically every spiral galaxy (Sancisi 1983, Briggs 1990) has been interpreted as a sign that a spiral galaxy accretes mass and angular momentum: the amount accreted is such that it changes completely orientation in a typical time-scale of a few Gyrs (Jiang \\& Binney, 1999).  During its life, a bar changes its pattern speed, and this could be a possible way to trace its age. In dissipationless simulations, the pattern speed decreases in a time-scale of a few Gyrs (Combes \\& Sanders 1981). This is also related to the length of the bar. The early bar instability involves central orbits, where the orbital precession rate is higher, then more and more orbits become trapped in the bar, at larger radii, and the bar lengthens and slows down, The transient spiral waves formed in the stellar component transfer the angular momentum outwards. This process is also accelerated by escaping chaotic orbits, around corotation (Pfenniger \\& Friedli 1991). If there is a massive spheroidal dark matter halo, concentrated enough to perturb the bar dynamics, the pattern speed of the bar could decrease even further, owing to the dynamical friction (Debattista \\& Sellwood 1998). According to the density of dark matter, the time scale could be estimated down to a few 10$^8$ yr. Since bars are observed most of the time to be fast rotators, this puts constraints on the density of dark matter in the central parts of galaxies. Fluctuations of the pattern speed can also be triggered by interacting companions (Gerin et al. 1990, Miwa \\& Noguchi 1998), but these remain transient. If gas accretion is able to rejuvenate the bars, it is important to know the implications on the pattern speed, that could be a tracer of the event. This is studied in detail in the presented simulations.  In Sect.~2 we briefly present the code and numerical methods used. In Sects.~3, 4 and 5 we present the results from our simulations. We give our conclusions in Sect.~6. ",
        "conclusions": "The numerical simulations presented in this work reveal how different can be the evolution of spiral galaxies with and without important gas accretion. While an isolated galaxy may spontaneously form a bar, consume its gas in stars, and when the stellar component is hot enough, see its bar weaken to a lens, and evolve definitely to an early-type system, a galaxy with significant gas accretion, can reverse this evolution, and experience several new bar episodes. The detailed processes are controlled by the strength of the bar and its gravity torques. The accreted gas may stay for a while in the outer parts of the disk, prevented to flow in by a strong bar; as soon as the bar weakens, the gas will progressively populate the disk, and form new stars there. This increases the self-gravity of the disk with respect to the stabilizing effect of the central bulge,  favoring a new bar instability. Under the influence of this new bar, the material inside its corotation is driven inwards, and through vertical resonance, thickens to enrich the bulge again. This regulates the bar strength.  In the course of the simulations, three to four bar episodes have been followed. An interesting discovery is that the pattern speed of the bar is increasing form one bar to the next. This can be considered as a rejuvenating process. Indeed, the normal evolution of a bar is to slow down, while trapping more and more material, and lengthening, so that its corotation radius increases (a bar cannot extend beyond its corotation, since periodic orbits are perpendicular here). When another bar forms, it is shorter than the previous one. Its pattern speed is higher, not only because there is then more mass contained in a given radius, due to mass accretion. The pattern speed is even higher than what would be expected from the square root of mass increase. This increase may be related to the fact that mass is more and more concentrated during this evolution, and the galaxy is shifting progressively to early-types, with massive bulges. The bar is then shorter with respect to the galaxy radius (as shown by Combes \\& Elmegreen 1993). Also, gas accretion has the effect of increasing the inner angular momentum per unit mass. The bar is shorter and rounder.  In the course of evolution with gas accretion, the radius of the galaxy expands, a phenomenon also controlled by the bar strength, which gravity torques can maintain or expel the gas outwards. One of the most important consequence of gas accretion, is to maintain an almost permanent (though recurrent) spiral structure through the disk, and also a significant star formation rate along a Hubble time. This solves the problem of explaining the almost constant star formation rate of spiral galaxies along their history (e.g. Kennicutt 1983), and also the presence until today of conspicuous spiral structure in most spiral disks.  The possibility to accrete gas from a direction highly inclined with respect to the galaxy plane introduces new processes that can be also favorable to bar reformation. The gas may not have to wait for a weakening of the previous bar to flow in, and the distribution is then asymmetric, with an m=2 extra contribution, which makes bars become much stronger.  One of the consequences of the evolution described in this work is that the bar pattern speed can no longer give the evolution state or \"age\" of a barred galaxy. Significant gas accretion can reverse the evolution and increase the pattern speed. Also, the argument that dynamical friction would have slowed down the bars along their evolution would have to be revised (e.g. Debattista \\& Sellwood 1998). Another consequence is that the morphology of a galaxy, according to its environment, and the availability to accrete cold gas, may oscillate from late to early-types three or four times, before inexorably shift to an early-type galaxy.  The maintenance of conspicuous spiral structure in most disk galaxies may mean that galaxies have sufficient gas to accrete in their outer parts, to constantly renew disk instability, even in the presence of a heating stellar disk and increasing bulge mass. The relative absence of bars in galaxies at high redshift (e.g. van den Bergh et al. 1996) could be interpreted as a very short life-time of the bar episode, in the presence of dominating gas fraction in very young galaxies. Indeed, the fraction of barred galaxies at a given epoch is interpreted, in this recurrent bar scenario, as the relative time spent by a galaxy in the bar episode relative to the un-barred one.  The present work has only considered general properties of gas accreting spiral galaxies, but many parameters (mass distribution, morphological types, geometry of the accretion, etc..) have to be explored before all consequences of accretion on galaxies can be drawn, which will be done in future work.  \\begin{appendix}"
    },
    "0206/astro-ph0206221_arXiv.txt": {
        "abstract": "We have analyzed the low mass X-ray binary (LMXB) candidates in a {\\it Chandra} observation of the giant elliptical galaxy NGC 4472. In a region observed by the Hubble Space Telescope (HST),  approximately 40\\% of the bright (L$_X$$\\gtrsim$10$^{37}$ ergs s$^{-1}$) LMXBs are associated with optically identified globular clusters (GC). This is significantly higher than the fraction of bright LMXBs in  Galactic GCs and confirms that GCs are the dominant sites of LMXB formation in early type galaxies. The $\\approx$4\\% of NGC 4472 GCs hosting bright LMXBs, on the other hand, is remarkably similar to the fraction of GCs with LMXBs in every other galaxy. Although statistical tests suggest that the luminosity of a cluster is an important driver of LMXB formation in GCs, this appears largely to be a consequence of the greater number of stars in bright clusters. The metallicity of GCs is a strong determinant of LMXB specific frequency, with  metal-rich clusters about 3 times more likely to host  LMXBs than metal-poor ones. There are weaker dependences on the size of a GC and its distance from the center of the galaxy. The X-ray luminosity does not depend significantly on the properties of the host GC. Furthermore, the spatial distribution and X-ray luminosity function of LMXBs within and outside GCs are indistinguishable. The X-ray luminosity function of {\\it both} GC-LMXBs and non-GC-LMXBs reveal a break at $\\approx$3$\\times$10$^{38}$ ergs s$^{-1}$ strongly suggesting that the brightest LMXBs are black hole accretors. ",
        "introduction": "Early X-ray studies with the {\\it Einstein} X-ray observatory revealed that elliptical  and S0 galaxies are significant sources of X-ray emission. While the bulk of the soft X-ray luminosity of X-ray bright galaxies can be attributed to kT $\\sim$ 1 keV thermal emission from hot (10$^7$ K) gas (e.g. Forman, Jones \\& Tucker 1985), the hard X-ray flux is roughly proportional to the optical luminosity, suggesting a low-mass X-ray binary origin (Trinchieri \\& Fabbiano 1985). With the advent of the {\\it Chandra} X-ray Observatory, it is now possible to resolve the point source X-ray populations likely to be associated with LMXBs in external galaxies such as NGC 4697 (Sarazin, Irwin \\& Bregman 2000) and NGC 1399 (Angelini, Loewenstein \\& Mushotzky 2001; hereafter A01), and definitively show that LMXBs make a significant contribution to the X-ray flux in typical ellipticals. Since the stellar populations in such galaxies are at least a few Gyrs old (e.g. Trager et al. 2000), contamination by high mass X-ray binaries is not an issue.  Globular clusters are especially fertile environments for LMXB formation.  Even though GCs account for $\\lesssim$0.1\\% of the stellar mass in the Galaxy, they harbor about 10\\% of the L$_X$$\\gtrsim$10$^{36}$ erg s$^{-1}$ LMXBs (e.g. Verbunt 2002), indicating  a probability of hosting a LMXB that is at least two orders of magnitude larger than for field stars. Recent {\\it Chandra} observations indicate that at least 20\\% of the LMXBs in NGC 4697 (Sarazin, Irwin, \\& Bregman 2001) and as much as 70\\% in NGC 1399 (A01) may reside in GCs, suggesting that GCs are even more dominant sites of LMXB formation in galaxies without recent star formation.  Dynamical processes, such as tidal capture of neutron stars  in close encounters with other cluster stars (Clark 1975; Fabian et al. 1975), and interactions between single stars and binaries (Hills 1976) in the high density environment of GC cores are some of the leading explanations for the preponderance of LMXBs in GCs. However, the very small number (13) of bright LMXBs in Galactic globulars (e.g. Liu, van Paradijs \\& van den Heuvel 2001; hereafter L01]) and other local group GCs has been the biggest impediment in testing these models in any detail. Thus, in order to isolate and understand the primary environmental factors driving LMXB formation in GCs it is imperative to study them in distant galaxies, especially globular cluster rich early types. In this Letter we analyze the LMXB population of NGC 4472, the giant elliptical in Virgo, and its connection to the GC system. ",
        "conclusions": "Even though the WFPC2 images cover only $\\approx$20\\% of the ACIS-S3 frame, 72 of the 144 LMXB candidates lie within this  region. We consider LMXBs to be associated with GCs if they are separated by less than 0.65$''$, where there is a natural break in the LMXB-GC angular separation. Thirty LMXBs satisfy this criterion. We note that increasing the matching distance to 1$''$ only adds 1 more source, while removing the color criterion adds 2 candidates with colors inconsistent with GCs. We choose not to relax the matching criterion in order to minimise the possibility of spurious matches.\t Thus, 40\\% of the L$_X$$>$10$^{37}$ ergs s$^{-1}$ LMXBs detected in NGC 4472 lie in GCs. This is significantly higher than the 5-15\\% figure in the Milky Way - based on the L01 catalog and Harris (1996) GC distances, with the variation due to how transient sources are counted - but smaller than the $\\approx$70\\% measured in NGC 1399 (A01).  The fraction of LMXBs in GCs and other associated numbers reported here represent ``snapshot\" values at the observational epoch. While it is possible that there are underlying populations of long-lived transients in the field or GCs, the instantaneous values are central to addressing scientifically interesting issues such as the importance of LMXBs to the X-ray emission from galaxies, their effect on X-ray background measurements and the significance of GCs to these values. Hence it appears that at present GCs are the dominant sites for LMXB formation in early type galaxies, in consonance with the suggestion of White, Sarazin \\& Kulkarni (2002).  Thus,  it is vitally important to study the GC-LMXB connection in these galaxies.   Although the fraction of LMXBs in GCs is different across galaxy types, the fraction of GCs hosting LMXBs is remarkably similar in all galaxies studied to date. Approximately 4\\% of NGC 4472 GCs host L$_X$$\\gtrsim$10$^{37}$ erg s$^{-1}$ LMXBs, compared to about 4\\% in NGC 1399 (A01),  $\\approx$2-3\\% in M31 (using DiStefano et al. 2002 and Barmby \\& Huchra 2001), and 1-4\\% in the Galaxy (using Liu et al. 2001 and Harris 1996). Thus the formation efficiency of LMXBs in GCs must be strongly driven by the properties of GCs, rather than that of the host galaxy.  \\subsection{Which Clusters Preferentially Form LMXBs?} The color distribution of  the NGC 4472 GCs is clearly bimodal (Fig. 2), with a population of blue, metal-poor GCs, and one of red, metal-rich GCs, as found in most other ellipticals (e.g. Zepf \\& Ashman 1993, KW01). KMM mixture modeling tests (Ashman, Bird \\& Zepf 1994) reveal a blue peak at V-I=0.98 and a red one at V-I=1.23, with a dividing color of V-I=1.10. One of the more striking results (Fig. 2) is that there are 3.3 times as many LMXBs in red GCs as there are in the blue. Even after accounting for the slightly larger number of red GCs, the overabundance factor is still $\\approx$2.7. Previous observations of Galactic GCs have also showed such a trend (e.g. Bellazzini et al. 1995), although small number statistics make it difficult to disentangle it from other GC properties such as distance from the center of the Galaxy.  The top panel of Fig. 2 also suggests a  tendency for more efficient LMXB formation in GCs in the inner region. The lower panel of Fig. 2 clearly shows that LMXBs are formed preferentially in the brightest GCs. This effect is actually even stronger than indicated when one considers that the faint end of the globular cluster luminosity function (GCLF) has not been corrected for incompleteness. Figure 2 also suggests that smaller GCs are favored LMXB hosts. Thus, each of these variables could provide the physical environment that promotes LMXB formation in GCs. In order to better understand the relative importance of each of these factors we turn to discriminant analysis (DA).  DA is used to weight and combine the discriminating variables in such a way that the differences between pre-defined groups are maximized (e.g. Antonello \\& Raffaelli 1983). Thus each data point is assigned a discriminant score of the form $F = w_1x_1 + w_2x_2 +  ... w_ix_i$ where F is the discriminant score, $w_i$ is the weighting coefficient for variable i, and $x_i$ is the i$^{th}$ discriminating variable, such that the distribution of discriminant scores of the pre-defined groups is maximally separated along the axis of this new composite variable. The absolute values of the standardized  coefficients,  $w_i$,  reveal the relative importance of the associated discriminating variables. In certain cases, where the discriminating variables may be correlated, the absolute value of the structure coefficients - which are the correlations of each variable with the discriminant function - may give better estimates of the significance of each of the variables.  Using the SPSSv10 package we performed DA on the LMXB and non-LMXB GC populations with V, V-I, r$_h$ and distance from the center of the galaxy as the variables. The standardized coefficients and structure coefficients (within brackets) for selected tests are presented in Table 1.  Since DA is a linear method, we tested the importance of GC size by using r$_h$, r$_h^3$ and log(r$_h$) in turn. We report only the r$_h$ analysis, which consistently assigns the most power to the size parameter. Two random variables, a Gaussian, and a uniform distribution are also used to gauge the significance of the results. Fig. 3 shows the success of DA in separating the two populations. DA of the entire sample clearly shows that the luminosity of a GC is  the most important factor that drives LMXB formation in GCs.  Color and distance appear to be equally important secondary factors that drive LMXB formation, with size providing a much weaker discriminant. We note however that the incompleteness of the faint end of the GCLF has a radial dependance that is both a function of the background light and differences in the exposure times of the WFPC2 images. Since the GCLF is constant at all radii (Kundu, Zepf \\& Ashman 2002) and the luminosity is known to be uncorrelated with color and r$_h$, restricting the sample to V$<$23.5 mag GCs, where the completeness is $\\approx$100\\%, provides a fairer statistical test for distance effects. This sample reveals that V, and V-I are the two most important discriminating variables, followed by smaller contributions from r$_h$ and distance. An independent statistical test using logistic regression analysis  reveals similar weights for each parameter, strengthening our confidence in the results.  The strong dependence of LMXB specific frequency on GC luminosity could    simply be a reflection of the larger number of stars in more luminous GCs, or it could possibly signal more efficient dynamical LMXB formation in massive GCs. The inset in Fig. 3 shows that the LMXB density per unit GC luminosity is roughly constant with V magnitude. Since GC luminosity is proportional to the total number of stars for the roughly constant M/L (and similar IMFs) in GCs , this suggests that the dependence of LMXB frequency on luminosity can largely be attributed to the greater number of stars in bright GCs.  It is also clear from our analysis that metallicity is a significant, strong, independent parameter. There are few theoretical explanations that can account for this. One possibility is the suggestion of Bellazzini et al (1995), that the larger radii of metal-rich stars promotes LMXB formation in GCs and facilitates mass transfer by Roche lobe overflow. While Grindlay (1987) suggested that a similar metallicity trend seen in Galactic GCs can be explained by a flatter initial mass function in metal-rich GCs, recent studies suggest that the observed differences in the present day mass function depend primarily on GC evolution, and not metallicity (e.g. Piotto \\& Zoccali 1999). While the rate of LMXB formation depends on the GC environment, our observations reveal no obvious trend of  L$_X$ with GC luminosity, color, size or distance. Further theoretical studies of the effect of GC environment on LMXB formation and evolution are required to understand the underlying physics.  \\subsection{LMXBs in Clusters vs. LMXBs in the Field} Are LMXBs in GCs different from those outside GCs? Were the non-GC-LMXBs in ellipticals formed in GCs and consequently ejected? In order to address some of these questions we compare the GC and non-GC populations of LMXBs. The mean luminosity of GC-LMXBs, $<L_X>=2.4\\pm0.6 \\times 10^{38}$ erg s$^{-1}$ is slightly higher than, but statistically indistinguishable from, the non-GC-LMXBs at $<L_X>= 1.8\\pm0.3 \\times 10^{38}$ erg s$^{-1}$.  Similarly, Kolmogorov-Smirnov tests revealed that the possibility that the GC and non-GC-LMXBs are drawn from the same L$_X$ or radial distance populations cannot be eliminated at greater than the 30\\% level. However, the exact distribution of each of these quantities in the two populations may hold vital clues about their origin. A single power-law fit  of the cumulative luminosity function (Fig. 4) of each of the LMXB categories can be eliminated at greater than 95\\% confidence level, while a broken power law of the form $N(>L) = N_{cut}(\\frac{L}{L_{cut}})^{\\alpha_{above/below}}$ fits each of the populations well. The non-GC-LMXBs and ``All\" source lists were corrected for background/foreground contamination, estimated from Mushotzky et al (2000), before fitting.  As is evident from Fig. 4, all populations show a break at L$_{cut}$$\\approx$3$\\times$10$^{38}$ ergs s$^{-1}$. A similar break is observed in other galaxies; Sarazin et al (2000) argue that the ``knee\" is at the typical Eddington luminosity for spherical accretion onto a 1.4 M$_\\odot$ neutron star, and may indicate the break between neutron star and black hole (BH) accretors. The luminous GC-LMXB population cannot be explained by multiple bright LMXBs in favorable GC environments as we see no trend of higher X-ray luminosity in preferred GC hosts. Also, the  ``knee\" in the non-GC-LMXB population of NGC 4472 cannot be caused by multiple LMXBs. Thus BH-LMXBs appear to be the most plausible explanation for the break based on present observations.     If non-GC-LMXBs were formed in clusters they can escape from GCs by one of two mechanisms, cluster destruction  - which is more efficient in the inner regions of galaxies (e.g. Vesperini 2000), and should reveal a surfeit of non-GC-LMXBs in the inner regions - or dynamical kicks that eject LMXBs from GCs. In the latter case, one might expect a spatial density distribution at least as extended as that of the GCs. If non-GC-LMXBs are formed in the field population, one would of course expect the density distribution to follow the light profile (which is steeper than the GC profile). The bottom panel of Fig. 4 suggests that the radial density distribution of GC-LMXBs and non-GC-LMXBs is roughly similar, and consistent with the GC population. At face value this suggests that dynamically ejected LMXBs from GCs may account for a significant fraction of non-GC-LMXBs.  Studies of the spatial properties of LMXBs over a larger radial range, especially the innermost regions of X-ray faint galaxies, are required to probe whether most non-GC-LMXBs in early type galaxies have a GC ancestry. Second epoch {\\it Chandra} observations will further allow discrimination between possible long-lived transient LMXBs associated with the field, and the more persistent sources that may be expected to form in GCs (Piro \\& Bildsten 2002).   SEZ and AK gratefully acknowledge support from NASA via the LTSA grant NAG5-11319. We thank Enrico Vesperini for many useful discussions."
    },
    "0206/astro-ph0206017_arXiv.txt": {
        "abstract": "We present the first complete 3-dimensional simulations of the core-collapse of a massive star from the onset of collapse to the resultant supernova explosion.  We compare the structure of the convective instabilities that occur in 3-dimensional models with those of past 2-dimensional simulations.  Although the convective instabilities are clearly 3-dimensional in nature, we find that both the size-scale of the flows and the net enhancement to neutrino heating does not differ greatly between 2- and 3-dimensional models.  The explosion energy, explosion timescale, and remnant mass does not differ by more than 10\\% between 2- and 3-dimensional simulations. ",
        "introduction": "Convective instabilities have been invoked to help drive core-collapse supernova explosions since Epstein (1979) first argued that negative lepton gradients would drive Ledoux convection in the core.  Epstein (1979) argued that this convection would increase the transport of energy out of the core and help facilitate a supernova explosion. Bruenn, Buchler, \\& Livio (1979) confirmed that this convection could indeed increase the neutrino luminosity and help drive a supernova explosion.  Considerable work studying convective instabilities, including multi-dimensional models (e.g. Buchler, Livio, \\& Colgate 1980), followed soon after.  Although entropy gradients caused by the shock were suggested during this time (see Bruenn, Buchler, \\& Livio 1979), Burrows (1987) first suggested that this entropy-driven convection could also boost the neutrino luminosity and help drive a supernova explosion.  The work of the past two decades has led to two separate convective regions: one within the extremely dense proto-neutron star core (see Keil, Janka, \\& M\\\"uller 1996 for a review) and the other in the region between the proto-neutron star and the accretion shock where the bounce stalled.  In this latter region, neutrino heating powers an unstable entropy gradient that drives convection (see Bethe 1990 for a review).  In the dense proto-neutron star, convection driven by lepton gradients (Epstein 1979, Keil et al. 1996), entropy gradients (Burrows 1987, Burrows \\& Lattimer 1988), and doubly diffusive (``salt-finger'') instabilities (Mayle \\& Wilson 1988) have all been invoked to increase the neutrino luminosity and hence, the neutrino heating.  In the neutrino heating region, entropy-driven convection helps to convert thermal energy gained from neutrino heating into kinetic energy, improving the over-all efficiency at which neutrinos from the core deposit energy into the outer layers of the star.  This latter convection has been studied in a number of 2-dimensional simulations over the last decade \\citep{Mil93,Heretal94,Bur95,Jan96, Mez98}.  This entropy driven convection occurs shortly after the collapse of the massive star.  When this core reaches nuclear densities and nuclear forces rapidly raise the pressure, its collapse halts, sending a bounce shock through the star.  This bounce shock stalls and leaves behind an unstable entropy profile that seeds convection in the region between the proto-neutron star and the edge of the stalled supernova shock.  Neutrinos leak out from the proto-neutron star and heat this region, continuing to drive this entropy-driven convection.  It is this convection that many groups now agree helps drive the supernova explosion \\citep{Heretal94,Bur95,Jan96}.  However, due to limitations in computer hardware and simulation software, all of the past work was limited to 2-dimensional simulations, leaving behind a number of unanswered questions.  Whether or not this increased efficiency is sufficient to drive a supernova explosion with the current supernova mechanisms is still a matter of debate: compare the explosions of \\citet{Heretal94} and \\citet{Bur95} to the fizzles of \\citet{Mez98}.  A key uncertainty in all of these simulations lies in the fact that the 2-dimensional simulations are being used to study an inherently 3-dimensional event in nature.  Some scientists have argued that nature will produce convective instabilities that are much different than what we see in the current 2-dimensional simulations.  In other convective problems (e.g. novae) it has been found that 2-dimensional models of these inherently 3-dimensional processes can lead to vastly incorrect answers (compare the differences between the 2- and 3-dimensional work of Kercek, Hillebrandt, \\& Truran 1998, 1999).  In this letter, we present the first complete 3-dimensional simulations of the evolution of a massive star from collapse to explosion, with particular emphasis on the differences between 2 and 3-dimensional models of the entropy-driven convection.  We follow these simulations until a strong supernova shock has been launched, and can hence see how these differences affect the final explosion energy, remnant mass, and nucleosynthetic yield of these supernovae. ",
        "conclusions": ""
    },
    "0206/astro-ph0206367_arXiv.txt": {
        "abstract": "We present spectroscopic observations of candidate F, G and K type stars in NGC 6633, an open cluster with a similar age to the Hyades. From the radial velocities and metal-line equivalent widths we identify 10 new cluster members including one short period binary system. Combining this survey with that of Jeffries (1997), we identify a total of 30 solar-type members. We have used the F and early G stars to spectroscopically estimate [Fe/H]\\,$=-0.096\\pm0.081$ for NGC 6633. When compared with iron abundances in other clusters, determined in a strictly comparable way, we can say with more precision that NGC 6633 has $(0.074\\pm0.041)$\\,dex less iron than the Pleiades and $(0.206\\pm0.040)$\\,dex less iron than the Hyades. A photometric estimate of the overall metallicity from the locus of cluster members in the {\\em B-V}, {\\em V-I}$_{\\rm c}$ plane, yields [M/H]\\,$=-0.04\\pm0.10$. A new estimate, based upon isochrones that are empirically tuned to fit the Pleiades, gives a distance modulus to NGC 6633 that is $2.41\\pm0.09$ larger than the Pleiades.  Lithium abundances have been estimated for the NGC 6633 members and compared with consistently determined Li abundances in other clusters. Several mid F stars in NGC 6633 show strong Li depletion at approximately the same effective temperature that this phenomenon is seen in the Hyades. At cooler temperatures the Li abundance patterns in several open clusters with similar ages (NGC 6633, Hyades, Praesepe and Coma Berenices) are remarkably similar, despite their differing [Fe/H].  There is however evidence that the late G and K stars of NGC 6633 have depleted less Li than their Hyades counterparts. This qualitatively agrees with models for pre-main sequence Li depletion that feature only convective mixing, but these models cannot simultaneously explain why these stars have in turn depleted Li by more than 1 dex compared with their ZAMS counterparts in the Pleiades.  Two explanations are put forward. The first is that elemental abundance ratios, particularly [O/Fe], may have non-solar values in NGC 6633 and would have to be higher than in either the Hyades or Pleiades. The second is that additional non-convective mixing, driven by angular momentum loss, causes additional photospheric Li depletion during the first few hundred Myr of main sequence evolution. ",
        "introduction": "Solar-type stars in open clusters are the obvious laboratories in which to study the evolution and timescales of a variety of physical phenomena (e.g.  magnetic activity, rotation, mixing).  Over the last couple of decades the Pleiades and Hyades, with ages of approximately 100\\,Myr and 700\\,Myr, have been the basis of much that has been deduced about the time-scales for the decline of rotation rates, X-ray activity and surface lithium abundances in solar type stars (e.g. Stern et al. 1992, 1995; Soderblom et al. 1993a,b; Thorburn et al. 1993; Stauffer et al. 1994; Krishnamurthi et al. 1997, 1998).  Consideration of these clusters alone, is not sufficient. Younger and older clusters need to be (and have been) studied of course, but observing clusters with similar ages to the Pleiades and Hyades is also important. Reasons include the possibilities: (a) that initial angular momentum or binary fractions are different from cluster to cluster, influencing their later behaviour; (b) that differing compositions or abundance ratios affect convection zone properties, which then feed in to the physical processes mentioned above.  An important illustration of this is the, as yet unexplained, different X-ray luminosity functions of solar-type stars in the Hyades and Praesepe, even though they share similar ages (Randich \\& Schmitt 1995).  In the last few years we have been adding to this database by studying the open cluster NGC 6633 ($=$ C 1825$+$065, Jeffries 1997; Briggs et al. 2000; Harmer et al. 2001).  The age of this cluster is found to be similar to the Hyades and Praesepe by a number of authors by looking at the main sequence turn-off and position of evolved stars in the Hertzsprung-Russell diagram (e.g. Harris 1976; Mermilliod 1981). However, the metallicity of NGC 6633 may be lower than either. Schmidt (1976), using $ubvy\\beta$ photometry, estimated a metallicity 0.2 dex lower than the Hyades and a distance of 348\\,pc.  Cameron (1985) gives [M/H] of $-0.13$, $+0.08$ and $+0.04$ for NGC 6633, the Hyades and Praesepe using {\\em UBV} photometry, and also finds distance and reddening estimates for NGC 6633 of 336\\,pc and $E(B-V)=0.17$. The Lynga (1987) catalogue uses weighted means from several different studies (see Janes, Tilley \\& Lynga 1988) to give [M/H] of $-0.11$, $+0.12$, $+0.07$ and ages of 630\\,Myr, 710\\,Myr, 830\\,Myr for NGC 6633, the Hyades and Praesepe respectively and we will adopt these ages in the rest of the paper. A distance of 312\\,pc and $E(B-V)=0.17$ is also quoted for NGC 6633 by Lynga (1987), although we have the means in this paper to make an independent distance estimate.  Spectroscopic estimates of the metallicity are currently rather crude.  Jeffries (1997) estimates [Fe/H] between $-0.1$ and $+0.05$ for a range of possible reddenings.  In summary, NGC 6633 likely provides a slightly lower metallicity analogue of the Hyades and Praesepe at a similar age.  Membership for solar-type stars in NGC 6633 can come via several routes. Sanders (1973) presents proper motions for 497 stars, complete to nearly $V=13$. Only a small fraction of these are classed as probable members and are primarily brighter stars. Proper motion appears to be a poor discriminator for stars with $V>12$, probably because the mean cluster peculiar tangential velocity is very small (about 1.5\\kms\\ with respect to the field average) and the fraction of contaminating background stars increases rapidly at fainter magnitudes. In some cases it is possible to rule out cluster membership on the basis of a large proper motion. Photometry can be used to select stars close to the ZAMS in colour-magnitude and/or colour-colour diagrams. Photoelectric {\\em UBV} photometry for 161 stars with $5.73<{\\em V}<15.11$ was presented by Hiltner, Iriarte \\& Johnson (1958). This survey seems (by comparison with Sanders' work) complete to $V=11$ but severely incomplete at fainter magnitudes. This work was extended by Jeffries (1997) using {\\em BVI} CCD photometry that was nearly complete to $V=20$.  Radial velocities were determined to $\\sim2$\\kms, for candidate cluster members with $10.9<V<15.1$, resulting in a refined list of 21 F to early K-type cluster members which shared a common radial velocity. Several likely short-period cluster binaries were identified on the basis of high proper motion membership probabilities but variable radial velocities.  Lithium is destroyed in cool star interiors as a consequence of encounters with protons at $(2-3)\\times10^{6}$\\,K. Convection and perhaps other mixing processes bring Li-depleted material to the surface, where the photospheric abundance can be measured using the \\lii\\ 6708\\AA\\ resonance doublet. The Li depletion pattern in NGC 6633 shows both similarities and differences to that in Hyades stars of equivalent temperature (Jeffries 1997). Depletion among the F stars was small and indistinguishable from the Hyades, though there was an absence of evidence for the ``Boesgaard gap'' of severely Li-depleted Hyades mid F stars -- a phenomenon thought to be driven by non-convective mixing processes (e.g. Boesgaard \\& Tripicco 1986; Boesgaard \\& Budge 1988). The G and early K-stars show tentative evidence for less Li depletion than the Hyades. This is expected from standard models that incorporate only convective mixing. Pre main sequence (PMS) Li depletion among stars with ultimate ZAMS temperatures of $>5000$\\,K should be strongly composition dependent -- lower metallicity stars have cooler convection zone bases and burn Li less efficiently for the same photospheric temperature. However, the same standard models cannot also explain why the G and K stars of NGC 6633 have depleted much more Li than their counterparts in the younger Pleiades, because little depletion is predicted on the main sequence. There is accumulating evidence that age, rather than composition is the primary determinant of the Li depletion suffered by a star of a given mass, pointing to roles for both additional mixing and perhaps a mechanism that inhibits strong PMS Li depletion amongst metal-rich stars (Jeffries \\& James 1999; Jeffries 2000; Ford et al. 2001; Barrado y Navascu\\'{e}s, Deliyannis \\& Stauffer 2001).   The status of NGC 6633 in testing these ideas is hampered both by small number statistics and uncertainty in the cluster metallicity compared with the better studied Hyades and Pleiades.  The purpose of this paper is to extend the study of Jeffries (1997) and define a larger sample of solar-type members of NGC 6633. This enlarged sample can then be used to study X-ray activity (see Briggs et al. 2001; Harmer et al. 2001), to continue the investigation of lithium depletion among the low-mass stars of NGC 6633 and to provide the first precise spectroscopic estimate of the cluster iron abundance.  In Section 2 we describe the photometric catalogue from which spectroscopic targets were selected. In Section 3 we discuss the spectroscopic observations and their analysis, including measurement of radial and rotational velocities.  Section 4 presents these results and combines them with those from Jeffries (1997). A revised membership list is constructed, we discuss the status of some individually peculiar stars and estimate to what extent our sample is complete or contaminated with non-members.  Section 5 presents new estimates, both spectroscopic and photometric, of the metallicity of NGC 6633 in comparison with the Hyades and Pleiades.  In Section 6 we determine the Li abundances of our cluster candidates and compare the Li depletion pattern of NGC 6633 with other clusters and theoretical models. The results are discussed in Section 7 and our conclusions presented in Section 8. ",
        "conclusions": "We have extended the spectroscopic survey of F--K type photometrically selected NGC 6633 candidates which was begun by Jeffries (1997). By considering the stellar radial velocities and metal line EWs we have found an additional 10 strong cluster candidates, including one new single-lined, short period, spectroscopic binary. Using uniform selection techniques we have added these stars to the data in Jeffries (1997) and arrived at a total of 30 likely cluster members with $0.39<${\\em B-V}$<1.15$ and spectral types from early F to early K. The mean heliocentric radial velocity of these stars (uncorrected for general relativistic effects) is $-28.2\\pm1.0$\\kms.  We have spectroscopically estimated the iron abundance of NGC 6633 using 10 single F and early G-type stars and the {\\sc atlas 9} atmospheres. We find [Fe/H]$=-0.096\\pm 0.081$, where the error includes (and is dominated) by an allowance for systematic $T_{\\rm eff}$ errors caused by choice of a colour-$T_{\\rm eff}$ error and uncertainties in the cluster reddening. When strict comparison is made with stars from the Pleiades and Hyades which have their [Fe/H] measured in an identical way, we find that [Fe/H]$_{\\rm NGC 6633}$ - [Fe/H]$_{\\rm Pleiades} = -0.074\\pm0.041$ and that [Fe/H]$_{\\rm NGC 6633}$ - [Fe/H]$_{\\rm Hyades} = -0.206\\pm0.040$.  A photometric estimate of the metallicity using the $B-V$ versus $V-I_{\\rm c}$ locus yields [M/H]\\,$=-0.04\\pm0.10$. Thus NGC 6633 appears to be a metal-poor (or at least iron-poor) version of the Hyades at a similar age. An estimate of the cluster distance is found from the colour-magnitude diagrams of the cluster members by using empirically tuned isochrone fits to the Pleiades ZAMS. We find that NGC 6633 has a distance modulus that is $2.41\\pm0.09$ larger than the Pleiades.  We have derived Li abundances using the \\lii\\ 6708\\AA\\ resonance doublet and compared these abundances to those estimated in a consistent fashion for other open clusters. We find that the Li depletion patterns among the F and early G stars of a group of similarly aged clusters (the Hyades, Praesepe, Coma Ber and NGC 6633) are almost identical, despite their differing [Fe/H].  We can confirm the presence of severely Li-depleted F stars at around 6600\\,K in NGC 6633, in close agreement with the ``Boesgaard gap'' already identified in the other three clusters.  We have shown that the Li abundance patterns in the late G and early K stars of NGC 6633 and the Hyades are different. There is now firm evidence that Li can still be detected in NGC 6633 stars with $T_{\\rm eff}\\simeq5200$\\,K at abundances nearly 1 dex higher than the upper limits found for Hyades stars at the same temperature. This difference is qualitatively in agreement with the expected dependence of PMS Li depletion on metallicity. However, this dependence is not clearly seen at higher temperatures and neither can PMS Li depletion explain why the G/K stars of NGC 6633 have less Li than their counterparts in the Pleiades which have similar or even higher [Fe/H].  We outline two scenarios that might explain these observations. One is that the elemental abundances in these clusters have non-solar ratios, altering the interior opacities and resulting in different amounts of Li depletion than would be predicted in models that assume all elemental abundances scale with iron. This would require [O/Fe] to be higher in NGC 6633 than either the Hyades or Pleiades. Alternatively, we propose that mixing between the convection zone base and regions hot enough to destroy Li is effective. This would be responsible for depleting Li in NGC 6633 from a level similar to, or higher than, that in the Pleiades during the first $\\sim500$\\,Myr of main sequence evolution."
    },
    "0206/astro-ph0206151_arXiv.txt": {
        "abstract": "The {\\it HETE-2} (hereafter \\HETE) French Gamma Telescope (FREGATE) and the Wide-field X-ray Monitor (WXM) instruments detected a short ($t_{50} = 360$ msec in the FREGATE 85-300 keV energy band), hard gamma-ray burst (GRB) that occurred at 1578.72 SOD (00:26:18.72 UT) on 31 May 2002.  The WXM flight localization software produced a valid location in spacecraft (relative) coordinates.  However, since no on-board real-time star camera aspect was available, an absolute localization could not be disseminated.  A preliminary localization was reported as a GCN Position Notice at 01:54:22 UT, 88 min after the burst.  Further ground analysis produced a refined localization, which can be expressed as a 90\\% confidence rectangle that is 67 arcminutes in RA and 43 arcminutes in Dec (90\\% confidence region), centered at RA = +15$^{\\rm h}$ 14$^{\\rm m}$ 45$^{\\rm s}$, Dec = -19$^\\circ$ 21\\arcmin 35\\arcsec (J2000).  An IPN localization  of the burst was disseminated 18 hours after the GRB (Hurley et al. 2002b).  A refined IPN localization was disseminated $\\approx$ 5 days after the burst.  This hexagonal-shaped localization error region is centered on RA = 15$^{\\rm h}$ 15$^{\\rm m}$ 03.57$^{\\rm s}$, -19$^\\circ$ 24\\arcmin 51.00\\arcsec (J2000), and has an area of $\\approx$ 22 square arcminutes (99.7\\% confidence region).  The prompt localization of this short, hard GRB by \\HETE and the anti-Sun pointing of the \\HETE instruments, coupled with the refinement of the localization by the IPN, has made possible rapid follow-up observations of the burst at radio, optical, and X-ray wavelengths. The time history of GRB020531 at high ($> 30$ keV) energies consists of a short, intense spike followed by a much less intense secondary peak.  Its time history is thus similar to that seen in many short, hard bursts.  Analysis of the FREGATE and WXM time histories gives durations for the burst of $t_{50}$ = 1.36 s in the WXM 2 - 25 keV energy range, and 1.10 s, 0.86 s, 0.62 s, and 0.36 s in the FREGATE 6-13, 14-30, 31-84, and 85-400 keV energy bands.  The duration of the burst thus increases with decreasing energy, which is similar to the behavior of long GRBs.  The photon number flux, photon energy flux, and energy fluence of the burst  in the 50-300 keV energy band in 1.25 seconds are 3.0 ph cm$^{-2}$ s$^{-1}$,  $6.4 \\times 10^{-7}$ erg cm$^{-2}$ s$^{-1}$, and $8.0 \\times 10^{-7}$ erg cm$^{-2}$, respectively.  The spectrum of the burst evolves from hard to soft, which is also similar to long GRBs.  These similarities to the properties of long GRBs, and other similarities previously known, suggest that short, hard GRBs are closely related to long GRBs. ",
        "introduction": "It has been known for nearly a decade that gamma-ray bursts (GRBs) appear to fall into two classes: short ($\\approx$ 0.2 sec), harder bursts, which comprise 20-25\\% of all bursts; and long ($\\approx$ 20 sec), softer bursts, which comprise 75-80\\% of the total (Hurley et al. 1992; Lamb, Graziani, and Smith 1993; Kouveliotou et al. 1993). % Norris, Scargle \\& Bonnell 2000).   The spectra of the two classes of bursts differ: the spectra of the bursts become softer as the bursts become longer (Dezalay et al. 1992; Kouveliotou et al. 1993; Dezalay et al. 1996). There is also evidence that the brightness distributions of the two classes of bursts differ (Graziani and Lamb 1994; Belli 1997; Tavani 1998); however, the difference in the brightness distributions can be explained by the difference in their durations, which causes the sampling distance for short bursts to be smaller than for the long bursts (Graziani and Lamb 1994).  In addition the $V/V_{\\rm max}$ values (Schmidt 2001) and the angular distributions (Kouveliotou et al. 1993) of the two classes appear to be identical.  Thanks to the rapid dissemination of accurate GRB localizations by \\BeppoSAX (Costa et al. 1997), much has been learned in the past five years about GRBs.  This has included the discoveries that GRBs have X-ray (Costa et al. 1997), optical (van Paradijs et al. 1997), and radio  (Frail et al. 1997) afterglows.  Redshifts and host galaxies are now known for more than two dozen GRBs (see, e.g., Lamb 2002). However, all of these discoveries relate to long GRBs.  In contrast, to date nothing is known about the distance to or the nature of the short GRBs, despite extensive efforts.  The Burst and Transient Source Experiment (BATSE) on the \\CGRO localized localized numerous short GRBs in near-real time.  Although many had large error boxes, one (trigger 6788) was localized to an error circle of $\\sim 30$ square degrees, which was searched for an optical counterpart within ~12 s to a magnitude of 14.98 (Kehoe et al. 2001).  The results were negative.  The Third Interplanetary Network (IPN) derived localizations for four short GRBs (000607, 001025B, 001204, and 010119) with delays of 15--65 hours.  But in three of these cases, the opportunity for follow-up observations was compromised either by the burst being close to the Sun (000607, $65^\\circ$) or close to the Galactic plane (001025B, $b \\approx 4^\\circ$; 010119, $b \\approx 5^\\circ$).  Only one burst (001204) was optimally placed on the sky for follow-up observations.  However, in this case the delay (65 hours) in deriving a localization for the burst hampered follow-up efforts.  Despite an accepted BeppoSAX ToO program, no X-ray  follow-up observations were possible because of Sun-angle or other operational constraints, except in the case of 001204.  However, the delay in deriving the localization of this burst made the success of any X-ray follow-up observation unlikely, and therefore none was carried out.  In this Letter we report the detection and prompt localization of a short, hard GRB by {\\it HETE-2} (hereafter \\HETE) (Ricker et al. 2002a,b).  On 31 May 2002 at 1578.73 SOD (05:15:50.56) UT on 31 May 2002 UT, the HETE-2 French Gamma Telescope (FREGATE) instrument (Atteia et al. 2002) and the Wide-field X-ray Monitor (WXM) instrument (Kawai et al. 2002) detected a bright, short (duration $\\sim 300$ msec in the FREGATE 30-400 keV energy band), hard GRB.  The prompt localization of this short, hard GRB by \\HETE and the anti-Sun pointing of the \\HETE instruments has made possible rapid follow-up observations of it at radio, optical, and X-ray wavelengths.  We also describe the properties of GRB020531 derived from observations of it using the FREGATE and WXM instruments on \\HETE, which provided unprecedented spectral and temporal coverage of this short, hard GRB. ",
        "conclusions": "\\label{conclusions}  \\HETE has detected and localized a short, hard GRB, establishing its capability to do so -- and demonstrating that the detection and localization of short, hard GRBs in the hard x-ray energy band is possible.  This has important implications for Swift and for other future GRB missions.  The prompt, precise localization of GRB020531 by \\HETE and the IPN have allowed rapid follow-up observations, which have placed much more severe limits on the brightness of any radio and optical afterglows from short, hard GRBs.  The complement of soft x-ray, hard x-ray, and gamma-ray instruments (SXC, WXM, and FREGATE) on \\HETE provides unprecedented  temporal and spectral coverage of short, hard GRBs.  With the currently projected long orbital lifetime (> 10 yrs) and excellent health of the \\HETE spacecraft and instruments, the results described for GRB020531 in this Letter demonstrate that \\HETE can continue to provide an unprecedented opportunity to study short, hard GRBs, and possibly to determine the distance to and the nature of these bursts."
    },
    "0206/astro-ph0206198_arXiv.txt": {
        "abstract": "It has been suggested that high velocity clouds may be distributed throughout the Local Group and are therefore not in general associated with the Milky Way galaxy.  With the aim of testing this hypothesis, we have made observations in the H$\\alpha$ line of high velocity clouds selected as the most likely candidates for being at larger than average distances.  We have found H$\\alpha$ emission from 4 out of 5 of the observed clouds, suggesting that the clouds under study are being illuminated by a Lyman continuum flux greater than that of the metagalactic ionizing radiation.  Therefore, it appears likely that these clouds are in the Galactic halo and not distributed throughout the Local Group. ",
        "introduction": "Our understanding of the nature of high-velocity clouds (HVCs), defined as interstellar clouds moving at velocities not compatible with a simple model of differential Galactic rotation, is severely limited by our lack of knowledge of their distances.  Except in a few isolated cases (see Wakker 2001), their distances are very poorly constrained.  Several authors (Blitz et al. 1999; Braun \\& Burton 1999) have suggested that a subclass of the HVCs are dispersed throughout the Local group of galaxies, the remnants of its formation. This would place them much farther away (100 kpc to 1000 kpc) than other models that place them in the Galactic halo at $\\lesssim$ 10 kpc distances (e.g. Oort 1970, Bregman 1980).  Given their angular size and neutral column density, the larger distances would make them very massive objects.  We have tested this hypothesis by measuring the H$\\alpha$ intensity toward a collection of HVCs whose properties open the possibility that they may be at greater than average distances.  If it is indeed true that these clouds lie at great distances from the Galaxy, their neutral gas should not be substantially ionized by the weak metagalactic ionizing flux.  Weymann et al. (2001) set a 2$\\sigma$ upper limit of 8 mR for the H$\\alpha$ intensity toward an intergalactic H~I cloud and inferred upper limits to the metagalactic ionizing flux, $\\Phi_{\\rm o}$, between 2.5 $\\times$ 10$^{3}$ cm$^{-2}$ s$^{-1}$ and 1.0 $\\times$ 10$^{4}$ cm$^{-2}$ s$^{-1}$ with a preferred value of 5 $\\times$ 10$^{3}$ cm$^{-2}$ s$^{-1}$, the range resulting from uncertainties in the cloud geometry (see also Shull et al. 1999 for a review of constraints on $\\Phi_{\\rm o}$).  However, if the cloud is in the vicinity of the Milky Way, a source of Lyman continuum photons, it will be more strongly ionized and therefore glowing in H$\\alpha$ (see Bland-Hawthorn \\& Maloney 1999).  In this case the observations are a measure of the escape fraction of ionizing photons from the Milky Way, which is important both in understanding the radiation transfer within the Galaxy and for understanding the impact of the Galaxy on its environment. ",
        "conclusions": "Of the five candidate Local Group clouds for which we obtained good measurements, there is clear H$\\alpha$ emission from four.  The intensities are around 0.1 R, typical of the intensities measured from the large HVC complexes thought to lie in the Galactic halo.  We conclude from this that it is unlikely that these clouds are at 100 kpc distances from our Galaxy, because the metagalactic ionizing flux level is below that needed to produce the observed H$\\alpha$ surface brightness of even the faintest of the detections.  If they are ``Local Group'' objects, then the metagalactic ionizing flux needs to be much higher than previously thought.  However, recent observations of the outer H~I disk of M 31 (Madsen et al. 2001) appear to rule out such a large ionizing flux within the Local Group.  It is more likely that these clouds are not at great distances, but are instead in the Galactic halo and being ionized by the Milky Way.  If this latter scenario is true, these observations support a picture where the distribution of ionizing flux percolating through the disk is somewhat patchy and the Lyman continuum flux, F$_{LC}$, incident onto a cloud from the Galactic plane averages F$_{LC}$ $\\simeq$ 2 $\\times$ 10$^5$ cm$^{-2}$ s$^{-1}$ (assuming 0.1 R as a representative H$\\alpha$ intensity; see Tufte et al. 1998).  The extreme faintness of H$\\alpha$ from HVC 394 place this measurement among the lowest upper limits to the metagalactic radiation field, comparable to the measurements of Madsen et al. 2001, Weymann et al. 2001, and Vogel et al 1995.  Since the H~I cloud is optically thick in the Lyman continuum and optically thin to H$\\alpha$ photons, each Lyman continuum photon incident on the cloud will ionize a hydrogen atom, and each hydrogen recombination will produce on average 0.46 H$\\alpha$ photons (Martin 1988; Pengally 1964; case B, T = 10$^4$ K).  If we assume that all of the H$\\alpha$ arises from gas photoionized by a uniform isotropic metagalactic ionizing flux, then the upper limit on I$_{\\alpha}$ provides direct constraints on this quantity.  Since the WHAM beam is easily contained within the H~I cloud (see beam F in Figure 4), the metagalactic ionizing flux generates observable H$\\alpha$ from the front and back faces of the cloud and we approximate the geometry as a plane slab viewed normally. Under these assumptions, the measured upper limit for HVC 394 of I$_\\alpha$ $<$ 0.01 R corresponds to an upper limit to the metagalactic ionizing flux $\\Phi_{\\rm o}$ $<$ 1.1 $\\times$ 10$^4$ photons cm$^{-2}$ s$^{-1}$, where $\\Phi_{\\rm o}$ is the incident one-sided ionizing flux defined by Vogel et al. 1995 and Madsen et al. 2001.  If other processes contribute to ionizing hydrogen in this cloud, then $\\Phi_{\\rm o}$ must be even lower.  Perhaps HVC 518 and HVC 444 are fainter in H$\\alpha$ because they are further away than HVC 532, HVC 486, and the M, A, and C complexes. However, with the exception of HVC 394, these results all suggest that the clouds are located in the halo rather than the intergalactic medium.  Corroborating evidence comes from the recent H$\\alpha$ observations of Weiner et al. (2001) of a different set of compact clouds, where they found intensities ranging from 0.04 R to 1.6 R, with no well-measured clouds showing non-detections.  The strength of these conclusions are currently limited by the relatively small number of clouds sampled, and further observations of the compact HVCs will be necessary to increase the statistics.  Also, observations in other emission lines such as [S~II], [N~II], and [O~III] should help to illuminate further the characteristics of the compact population of HVCs and their relationship to the major complexes.  We are grateful to Dr. Leo Blitz for constant encouragement and support, and for his help in selecting HVCs to observe.  S. L. T. and J. D. W. acknowledge support from Research Corporation through a Cottrell College Science Award.  G. J. M., L. M. H., and R. J. R. acknowledge support from the National Science Foundation through grant AST 96-19424.   \\pagebreak"
    },
    "0206/astro-ph0206401_arXiv.txt": {
        "abstract": "{ Smoothed Particle Hydrodynamics is reformulated in terms of the convolution of the original hydrodynamics equations, and the new evolution equations for the particles are derived. The same evolution equation of motion is also derived using a new action principle. The force acting on each particle is determined by solving the Riemann problem. The use of the Riemann Solver strengthens the method, making it accurate for the study of phenomena with strong shocks. The prescription for the variable smoothing length is shown. These techniques are implemented in strict conservation form. The results of a few test problems are also shown. }  \\begin{article} ",
        "introduction": "Smoothed particle hydrodynamics (SPH, \\cite{Lucy:SPH,GM:SPH}) is a fully Lagrangian particle method to describe the hydrodynamical phenomena. The Lagrangian particle methods are especially suited to hydrodynamical problems that have large empty regions and moving boundaries. Those problems naturally arise in engineering science as well as geophysics and astrophysics. A variety of astrophysical problems have been studied by SPH because of its simplicity in programming the two- and three-dimensional codes and its versatility of incorporating the various physical effects such as self-gravity, radiative cooling, and chemical reactions. A broad discussion of the method can be found in a review by Monaghan \\cite{Monaghan:ARAA}. However, the inconsistency of the method is emphasized by Dilts \\cite{Dilts:1999,Dilts:2000} who modified the method by means of the moving-least-square basis functions to obtain accurate derivatives regardless of the positions of the SPH particles. Another concern of the method is its poor description of the strong shocks. In the two- or three-dimensional calculation of colliding gases, particles often penetrate into the opposite side. This unphysical effect can be partially eliminated by the so-called XSPH prescription \\cite{Monaghan:1989} that does not introduce the (required) additional dissipation but results in the additional dispersion of the waves. Therefore it is very desirable to construct a method that is simple and able to describe the strong shock phenomena accurately.  In this paper, a new method for handling shocks in particle hydrodynamics is constructed. The force acting on each particle is determined by solving the Riemann problem (RP). This use of the so-called Riemann Solver is introduced as an simple analogy of the grid-based method \\cite{Godunov:1959}. The previous attempts to introduce Riemann Solver into the particle method failed to give the method an exact conservation form \\cite{SI:GPH,SI:IAU95}. This paper describes how to include the exact Riemann Solver into the strictly conservative particle method.  The alternative approach to including small but sufficient dissipation into the numerical solution is to find a good limiter to switch a dissipative method and a less-dissipative method \\cite{FulkQuinn:1995,MorrisMonaghan:1997}. In principle, however, the switch is always an option to any numerical scheme including the present one, and its discussion is beyond the scope of this paper.  Section 2 provides a description of the method where we derive the evolution equations for the particles in terms of the convolution of the original hydrodynamic equations with the so-called kernel function. The same evolution equations are derived from an action principle which is different from the previous ones \\cite{GM:1982, Nelson:1994}, which sheds light on the ``hidden'' approximation in the expression for the velocity field in SPH formalism. The detailed explanation for the implementation is described in Section 3. The usage of the Riemann Solver is analogous to the grid-based second-order Godunov scheme \\cite{vL:1979}. A variable smoothing length is also considered. Numerical examples involving strong shocks are presented in Section 4. Section \\ref{Sec:Summary} is for summary. ",
        "conclusions": "\\label{Sec:Summary} Smoothed Particle Hydrodynamics is reformulated by the formal convolution of the original hydrodynamics equations, and by a new action principle, in which the second order (in $h$) approximation is used for the kinetic term of the Lagrangian function. The force acting on each particle is determined by solving the Riemann problem for each particle pair. The prescription for variable smoothing length is also shown. These techniques are implemented in the strict conservation form. Numerical examples involving an extremely strong shock are shown. The other test calculations will be described elsewhere.  Although the method with a spatially constant smoothing length is formulated in a rigorous manner, the method with the variable smoothing length is based on a crude approximation (Eqs. [\\ref{eq:EoM6.5}],[\\ref{eq:EoE6.5}]). A more refined technique for the variable smoothing length is to be studied. A better approximation than the cubic spline interpolation in the numerical convolution is required to eliminate the ``wiggle'' at the contact discontinuity in the blast wave problems (see Section \\ref{sec:BlastWave}).   This paper presented a piece of concepts that are not discussed in detail in the literature. Those are the convolution of the original fluid equation (Eqs. [\\ref{eq:EocEoM}],[\\ref{eq:EocEoE}]), the definition and approximation for the velocity field (Eqs. [\\ref{eq:Dov}], [\\ref{eq:Aov}]), and the modification of the force due to dissipation (Eqs. [\\ref{eq:EoM5}], [\\ref{eq:EoE5}]). Incorporating these concept, the final evolution equation was cast into a form similar to the standard SPH equation. Therefore those concepts may enable the rigorous examination of the accuracy and stability of the method and may enable further modification.  The numerical examples in Section \\ref{sec:example} show that the present particle method based on Riemann Solver can handle severe problems with strong shocks, those might include the description of explosion/implosion and supersonic jet phenomena. In this respect, further modification of the SPH method in modeling relativistic flows is promising with the help of relativistic Riemann Solver \\cite{MartiMuller:1996}. In adittion, the Lagrangian particle methods have advantage over Eulerian grid-based methods, in describing chemically reacting (multi)fluid and radiatively heating/cooling fluid \\cite{SI:NAP}. This is because we can simply assign the chemical composition and entropy to each particle as a fluid element. This direction has a wide area of applications.    \\begin{acknowledgment} The author thanks the anonymous referee for valuable comments. The author also thanks Toru Tsuribe, Yusuke Imaeda, and Shoken M. Miyama for useful discussions. \\end{acknowledgment}"
    },
    "0206/astro-ph0206292_arXiv.txt": {
        "abstract": "The abundance of clusters at the present epoch and weak gravitational lensing shear both constrain roughly the same combination of the power spectrum normalization $\\sigma_8$ and matter energy density $\\Omega_M$. The cluster constraint further depends on the normalization of the mass-temperature relation. Therefore, combining the weak lensing and cluster abundance data can be used to accurately calibrate the mass-temperature relation.  We discuss this approach and illustrate it using data from recent surveys. ",
        "introduction": "The number density of galaxy clusters as a function of their mass, the mass function, and its evolution can provide a powerful probe of models of large-scale structure. Historically the most important constraint coming from the present day abundance of rich clusters has been the normalization of the {\\it linear theory\\/} power spectrum of mass density perturbations (e.g.\\ \\cite{Evr89,FWED,BonMye91,HA91,Lil,OukBla,BahCen,WEF,VL96,VL98,Henry}). The normalization is typically quoted in terms of $\\sigma_8$, the rms density contrast on scales $8\\,h^{-1}\\,$Mpc, with the abundance constraint forcing models to a thin region in the $\\Omega_M$-$\\sigma_8$ plane.  Since the mass, suitably defined, of a cluster is not directly observable, one typically measures the abundance of clusters as a function of some other parameter which is used as a proxy for mass. Several options exist, but much attention has been focused recently on the X-ray temperature. Cosmological N-body simulations and observations suggest that X-ray temperature and mass are strongly correlated with little scatter (\\cite{EMN,BryNor,ENF,HorMS,NevMF}). How well simulations agree with observational results is far from clear, and several issues need to be resolved.  On the simulation side there are the usual issues of numerical resolution and difficulties with including all of the relevant physics.  On the observational side instrumental effects can be important (especially for the older generation of X-ray facilities) in addition to the worrying lack of a method for estimating ``the mass''. In this respect it is worth noting that there are numerous differing definitions of which ``M'' and ``T'' are to be related in the M--T relation (\\cite{White_mass})!  With current samples the {\\it dominant\\/} uncertainty in the normalization in fact comes from the normalization of the M--T relation (\\cite{ECF96,VL96,DV99,Henry,PSW,Seljak}). Or phrased another way, the cluster abundance is a sensitive probe of the normalization of the M--T relation.  The abundance of clusters is, of course, not the only way to constrain the cosmological parameters. In this regard it is interesting to note that weak gravitational lensing provides a constraint on a very similar combination of $\\Omega_M$ and $\\sigma_8$. Therefore, the two constraints can be combined to check for consistency of our cosmological model, to provide a normalization for the M--T relation, to probe systematics in either method and/or to measure other parameters not as yet included in the standard treatments.  While the cluster constraint comes primarily from scales of about $R=10\\,h^{-1}\\,$Mpc, current weak lensing surveys constrain somewhat smaller scales.  These surveys probe scales between roughly 1 and 10 arcmin, which for source galaxies located at $z\\simeq 1$ in a $\\Lambda$CDM cosmology corresponds to $0.7\\,h^{-1}\\,{\\rm Mpc}<R<7\\,h^{-1}\\,{\\rm Mpc}$.  Therefore, weak lensing probes slightly smaller scales than clusters.  As lensing surveys push to larger scales the overlap will become even better.  In this paper we argue that a natural application of combining the cluster abundance and weak lensing constraints is to calibrate the M--T relation for galaxy clusters (see also \\cite{HuKrav}).  In Sec.~\\ref{sec:MT} we define the M--T relation and derive how cluster abundance constraints depend on $\\Omega_M$ and $\\sigma_8$.  In Sec.~\\ref{sec:WL_clus} we illustrate how combining the two constraints can fix the normalization of the M--T relation using two recently obtained data sets.  Finally, in Sec.~\\ref{sec:concl} we discuss this approach further. ",
        "conclusions": "\\label{sec:concl}  There has been a lot of discussion recently regarding the value of cluster normalization $\\sigma_8$.  While the ``old'' results favor $\\sigma_8\\sim 1$ (\\cite{VL98,PSW} and references therein), several new cluster abundance analyses favor a significantly lower normalization (\\cite{ReiBoh,Bor01,Viana_Nichol,Seljak,Ikebe,SDSS}).  The lower normalization is also favored by the combined analysis of 2dF Galaxy Redshift Survey and CMB data (\\cite{2dF}).  On the other hand, recent weak lensing results (\\cite{vW02,Bacon,Refregier,Hoekstra}) tend to favor a higher value of $\\sigma_8$. The cause is of this discrepancy between various measurements has not been identified yet; one candidate is larger than anticipated systematic errors in one or both methods.  Another possibility is the bias in the relation between the mass and the observable quantity --- temperature or luminosity --- used to construct the abundance of clusters.  \\vspace{0.2cm} \\begin{figure}[!ht] \\epsfig{file=f2.eps, height=3in, width=3.5in} \\caption{Estimates of the M--T normalization $T_*$ collected from the literature. The red (light) points are estimates from simulations, while the black (dark) points are from the observations. Points with no error bars had none quoted. The shaded region is roughly our favored range of values of $T_*$.} \\label{fig:MT_lit} \\end{figure}  The cluster abundance constraint on $\\sigma_8$ crucially depends on the M--T normalization $T_*$.  Figure~\\ref{fig:MT_lit} summarizes the current status of our knowledge of $T_*$. It shows seven determinations from N-body simulations and three from direct observations, as compiled in Pierpaoli, Scott \\& White (2002) and Muanwong et al.\\ (2002). The shaded region is roughly our favored range of values of $T_*$. Points without error bars had none quoted, and the three observed values of $T_*$ assumed the isothermal-$\\beta$ model. The measurement due to Muanwong et al.\\ corresponds to their ``radiative'' and ``preheating'' cases that are cooling-flow corrected, while the value due to Pierpaoli, Scott \\& White is an average over the simulations.  The large discrepancy between the different measurements is apparent, and it also appears that the observed values are systematically higher than the ones obtained from simulations (see Muanwong et al.~2002 for further discussion).  We argue here that the cluster abundance -- weak lensing complementarity can be used to cross-check the M--T relation. By combining recent weak lensing constraints from Bacon et al.\\ and the HiFluGCS cluster sample of Reiprich \\& B{\\\" o}hringer, we have demonstrated the utility of this method. While potential systematic errors in both data sets are still a concern, the example we used prefers relatively low values of the M--T normalization ($T_*\\lesssim 1.7\\,$keV). We conclude that future weak lensing surveys (Vista, LSST, SNAP) combined with new cluster data from Chandra and XMM-Newton observations will provide a strong probe of the M-T relation."
    },
    "0206/astro-ph0206036_arXiv.txt": {
        "abstract": "According to high resolution cold dark matter (CDM) simulations, large virialized halos are formed through the constant merging of smaller halos formed at earlier times. In particular, the halo of our Galaxy may have hundreds of dark matter clumps. The annihilation of dark matter particles such as the neutralino in these clumps generates $\\gamma$-ray fluxes that can potentially be detected by future experiments such as GLAST. We find that, depending on the parameters of the clump density profile and on the distribution of clumps in the Galactic halo, the contribution to the diffuse $\\gamma$-ray background from clumps can constrain the properties of neutralinos such as the mass and annihilation cross section.  We model the density profile of clumps by three representative  dark matter profiles: singular isothermal spheres (SIS), Moore profiles, and Navarro, Frenk and White (NFW) density profiles and calculate  the spectrum and angular distribution in the sky of the $\\gamma$-ray  flux due to neutralino annihilation in the clumpy halo of the Galaxy.  The calculations are carried out in the context of two different scenarios for the distribution of clumps in the Galaxy and their concentrations, which  result in very different conclusions. ",
        "introduction": "Most of the matter in the universe has yet to be observed in any frequency band, thus the name, dark matter (DM). The evidence for the predominance of  dark over visible matter comes mainly from the gravitational effects of the dark  matter component. However, gravitational studies have been unable to shed light on the nature of the dark matter. Big bang nucleosynthesis constrains  most of the dark matter to be of non-baryonic origin. This has encouraged the study of plausible  new particle candidates for the dark matter. The best motivated among these candidates is the lightest supersymmetric particle, named the neutralino, $\\chi$. Given the requirements of neutralino production in the early universe, it is possible to study the phenomenology of such dark matter candidates in detail (see, e.g., Jungman, Kamionkowsky, and Griest 1996).  In particular, the annihilation of  neutralinos has often been considered  a potential source of detectable secondaries: high energy particles and  electromagnetic radiation. In this sense, dark matter can be visible  through the radiation caused by the annihilation secondaries (Berezinsky, Gurevich and Zybin 1992; Berezinsky, Bottino and Mignola 1994; Gondolo and Silk 1999; Bergstr\\\"om et al. 1999; Gondolo 2000; Calc\\'{a}neo--Rold\\'{a}n and Moore 2000; Bertone, Sigl and Silk 2001; Bergstr\\\"om, Edsj\\\"o and Ullio 2001; Blasi, Olinto and Tyler 2002; Tyler 2002; Ullio et al. 2002; Tasitsiomi and Olinto 2002; Cesarini et al. 2003).   Recent advances in cold dark matter (CDM) simulations have shown that the large scale structure of the universe can be explained in terms of a hierarchical scenario in which large halos of dark matter are generated by the continuous merging of smaller halos (Ghigna et al. 1998; Moore et al. 1999; Klypin et al. 1999). In this picture, a dark matter halo is the superposition of a smooth component on a scale comparable with the virial radius of the forming  structure and a clumped structure made of thousands of smaller mass halos. As we discuss below, this scenario can have important consequences for the detection of signals from neutralino-antineutralino ($\\chi\\bar{\\chi}$) annihilation.   CDM simulations also show that most halos are well described by a density distribution with cusps at the center of each halo.  The exact shape of the central cusp is still a matter of debate. Most recent simulations favor profiles with density cusps varying from the  Moore et al. profile (Moore et al. 1999) where $\\rho_{DM} (r \\to 0)\\sim r^{-1.5}$ to a Navarro, Frenk, and White (NFW) profile (Navarro, Frenk and White 1996) where $\\rho_{DM} (r \\to 0) \\sim r^{-1}$ (see also Power et al. 2002). The more concentrated the central cusp, the easier to detect annihilation products since the annihilation rate is proportional to the square of the dark matter density. The existence of cuspy halos is still unclear since observations of galaxy rotation curves provide no evidence  of central cusps (see, e.g., Salucci 2001). However, the survival of  cusps in galactic centers is highly dependent on the galaxy's merger  history in particular on the history of formation of galactic center  black holes (Merrit and Milosavljevic 2002).  The central regions of small mass dark matter halos (i.e., DM clumps) are less affected by the dynamics of baryonic matter and less likely to have black hole mergers at their centers. Thus, a cuspy profile may well describe the density of DM clumps.  Another important piece of information is embedded in the spatial distribution and survival history of clumps on their way to the central part of the host galactic halo. Tidal interactions significantly affect clump distribution and structure and these effects are only beginning to be resolved in CDM simulations for relatively large mass clumps (Stoehr et al. 2002). Given the lack of a clear understanding of present distribution of DM clumps in the Galaxy halo, we consider two  different scenarios that bracket the present uncertainties. For simplicity we call these two approaches type I and type II. We show that the present uncertainties lead to large variations in the resulting gamma-ray signal.  In the type I scenario, the clump position in the host Galaxy determines its external radius. In the case of NFW and Moore profiles, the location of the core is then determined by assuming a fixed fraction of this external radius. In the type II scenario, the concentration parameter, defined as the ratio of the core radius to virial radius of a clump, is modeled as in Wechsler et al. (2001) for clumps for a given mass. We show that in this scenario the clumps are  gradually destroyed on their way to the center of the Galaxy, such that the inner part of the Galactic halo is depleted of clumps. This second scenario seems compatible  with numerical simulations (Power et al. 2002, Stoehr et al. 2002). The uncertainties on clump distribution and concentration bracketed by these two scenarios strongly affect the resulting $\\gamma$-ray emission from the annihilation of neutralinos: while in the type I scenario, the highly concentrated  clumps produce strong $\\gamma$-ray emission, in the  type II scenario the low concentration clumps imply $\\gamma$-ray fluxes several orders of magnitude smaller.   The paper is structured as follows: in section 2 we introduce the clump density profiles; in section 3 we discuss the distribution of clumps in the Galaxy and the calculations of the gamma ray flux; in section 4 we present our results, consisting of the spectra and anisotropies of the $\\gamma$-ray emission in the type I and II scenarios, for different shapes of the clumps density profiles. We conclude in section 5. ",
        "conclusions": "The annihilation of dark matter in the halo of our galaxy generates a diffuse background of gamma-rays. The dark matter substructure of the galactic halo can be a dominant component of the diffuse gamma-ray flux depending on the concentration and location of DM clumps in the inner regions of the Galactic halo.  In order to bracket the range of possible fluxes, we considered two extreme scenarios, that we named type I and type II. The first corresponds to  extremely concentrated clumps present everywhere in the Galaxy halo, while in the second scenario the clumps are much less concentrated and are completely destroyed by tidal effects in the inner 10 kpc of the Galaxy. While the type I scenario allows one to put very strong constraints on the properties of neutralinos and on the density profile inside clumps, the type II scenario generates  fluxes of diffuse gamma-rays which are  barely detectable, with the exception  of the case in which the density profile of the clumps is modeled as a SIS sphere.  For the type I scenario, most of the parameter space of neutralino dark matter is ruled out if the density profile of dark matter clumps is in the form of a SIS sphere or a Moore profile. Weaker bounds can be imposed on the neutralino parameter space in the case of NFW density profile.  Recent N-body simulations have started to gather evidence of the destruction of clumps in the inner regions of dark matter halos, which favor type II like scenarios. On the other hand, numerical simulations only resolve relatively high mass clumps and the distribution of more concentrated smaller clumps is yet to be understood, although there is some recent progress in this direction (Berezinsky, Dokuchaev and Eroshenko 2003). Finally,  one should keep in mind the possibility of detecting isolated nearby clumps with Atmospheric Cherenkov telescopes (Tasitsiomi and Olinto 2002), which  may be more promising than looking at diffuse fluxes."
    },
    "0206/astro-ph0206200_arXiv.txt": {
        "abstract": "In this paper (paper I) we present optical long-slit spectroscopy and imaging of the protoplanetary nebula \\objectname{CRL 618}. The optical lobes of \\crl\\ consist of shock-excited gas, which emits many recombination and forbidden lines, and dust, which scatters light from the innermost regions. From the analysis of the scattered H$\\alpha$ emission, we derive a nebular inclination of $i$=24\\degr$\\pm$6\\degr. The spectrum of the innermost part of the east lobe (visible as a bright, compact nebulosity close to the star in the H$\\alpha$ $HST$ image) is remarkably different from that of the shocked lobes but similar to that of the inner \\ion{H}{2} region, suggesting that this region represents the outermost parts of the latter. We find a non-linear radial variation of the gas velocity along the lobes. The largest projected LSR velocities ($\\sim$\\,80\\,\\kms) are measured at the tips of the lobes, where the direct images show the presence of compact bow-shaped structures. The velocity of the shocks in \\crl\\ is in the range $\\sim$75-200\\,\\kms, as derived from diagnostic line ratios and line profiles. We report a brightening (weakening) of [\\ion{O}{3}]$\\lambda$5007\\AA\\ ([\\ion{O}{1}]$\\lambda$6300\\AA) over the last $\\sim$\\,10 years that may indicate a recent increase in the speed of the exciting shocks. From the analysis of the spatial variation of the nebular extinction, we find a large density contrast between the material inside the lobes and beyond them: the optical lobes seem to be `cavities' excavated in the AGB envelope by interaction with a more tenuous post-AGB wind. The electron density, with a mean value $n_{\\rm e}$\\,$\\sim$5$\\times$10$^3$-10$^4$\\,\\cm3, shows significant fluctuations but no systematic decrease along the lobes, in agreement with most line emission arising in a thin shell of shocked material (the lobe walls) rather than in the post-AGB wind filling the interior of the lobes. The masses of atomic and ionized gas, respectively, in the east (west) lobe are $>$1.3$\\times$10$^{-4}$\\,\\ms\\ ($>$7$\\times$10$^{-5}$\\,\\ms) and $\\sim$6$\\times$10$^{-5}$\\,\\ms\\ ($\\sim$4$\\times$10$^{-5}$\\,\\ms). The shocks in \\crl\\ are in a radiative regime and may lead in the future to the evolution of the optically-emitting lobes into a fast, bipolar molecular outflow. The time required by the dense, shocked gas to cool down significantly is $\\lsim$\\,2\\,yr, which is substantially lower than the kinematical age of the lobes ($\\lsim$\\,180\\,yr). This result suggests that a fast wind is currently active in \\crl\\ and keeps shocking the circumstellar material. ",
        "introduction": "\\label{intro} The physical mechanisms responsible for the onset of bipolarity and polar acceleration in planetary nebulae (PNe) are already active in the first stages of the evolution beyond the Asymptotic Giant Branch (AGB), i.e$.$ in proto-Planetary Nebulae (PPNe, also called post-AGB objects). Therefore, PPNe and young PNe hold the key for understanding the complex and fast ($\\sim$\\,10$^3$\\,yr) nebular evolution from the AGB towards the PN phase. Such evolution is believed (by an increasing number of astronomers) to be governed by the interaction between fast, collimated winds or jets, ejected in the late-AGB or early post-AGB phase, and the spherical and slowly expanding circumstellar envelope (CSE) resulted from the star mass-loss process during the AGB \\citep[see][]{sah98,kas00}.   \\objectname{CRL 618} (= \\objectname{RAFGL 618} = \\objectname{IRAS 04395+3601} = \\objectname{Westbrook Nebula}) is a well studied PPN which has very recently started its post-AGB journey \\citep[only $\\sim$\\,200\\,yr ago, e.g.][]{kwo84} and is rapidly evolving towards the PN stage.  Most of the circumstellar matter in \\crl\\ is still in the form of molecular gas. This component, with a total mass of $\\approx$\\,1.5\\ms, consists of: (1) a spherical and extended ($\\gsim$\\,20\\arcsec) envelope expanding at 17.5\\,\\kms\\ \\citep{kna85,buj88,buj01,bac88,haj96,phi92}; and (2) an inner, compact bipolar outflow moving away from the star at velocities $\\gsim$\\,200\\,\\kms\\ \\citep{cer89,gam89,mei98,ner92}. The outer, and slowly expanding component is interpreted as the result from the mass-loss process of the central star during the AGB, which took place at a rate of \\mloss\\,=\\,few$\\times$\\,10$^{-5}$-10$^{-4}$\\,\\my. The fast bipolar outflow, with a mass of $\\sim$10$^{-2}$\\ms, is believed to be the result of the interaction between a fast, collimated post-AGB wind and the spherical AGB CSE (see references above).    The high-excitation nebula (atomic and ionized gas) is composed of: (1) a compact \\ion{H}{2} region, visible through radio-continuum emission, elongated in the E-W direction with an angular size of 0\\farcs2-0\\farcs4 that is increasing with time \\citep{kwo81,mar93};  and (2) multiple lobes with shock-excited gas which produces recombination and forbidden line emission in the optical \\citep[e.g.][]{goo91,tra93,kel92,tra00}. From previous spectroscopic data, the inner \\ion{H}{2} region and the lobes are known to be expanding with velocities of $\\sim$\\,20\\,\\kms\\ and up to $\\sim$\\,120\\,\\kms, respectively \\citep{mar88,car82}. The analysis of different optical line ratios indicates that a relatively large range of temperatures ($\\approx$\\,10,000 to 25,000\\,K) and densities ($\\approx$\\,1600 to 8000\\,\\cm3) are present in the lobes \\citep{kel92}.  The optical spectrum of \\crl\\ also shows a weak, red continuum which is the stellar light reflected by the nebular dust. From spectropolarimetric observations it is known that a fraction of the flux of the Balmer lines is also scattered light originally arising from the inner, compact \\ion{H}{2} region \\citep{sch81,tra93}. The polarization of the forbidden lines is negligible, indicating that they are almost entirely produced by the shock-excited gas in the lobes with a small or insignificant contribution by scattered photons from the \\ion{H}{2} region (the high density in this region, $\\sim$10$^6$\\,\\cm3, produces collisional de-excitation of most forbidden lines).  The central star of CRL\\,618 has been classified as B0 based on the shape of the dereddened visual continuum \\citep{sch81} and on the weak [\\ion{O}{3}] line emission from the inner, compact \\ion{H}{2} region \\citep{kal78}. The luminosity of CRL\\,618, obtained by integrating the observed IRAS fluxes, is $L=1.22 \\times 10^4$\\ls$[D/{\\rm kpc}]^2$ \\citep{goo91}. Based on the $L$ values predicted by evolutionary models ($\\sim$\\,10$^4$\\ls) and the typical scale height of PN ($\\sim$\\,120\\,pc) these authors calculate a distance to the source of 0.9\\,kpc, which we also assume in this paper.  In this paper (paper I of a series of two) we report optical imaging and long-slit spectroscopy of \\crl. Observational techniques and data reduction are described in Sect$.$ \\ref{obs}. In Sections \\ref{morpho} and \\ref{sect4} we present our observations, including a brief description of the nebular morphology and main characteristics of the optical spectrum. In Sect$.$ \\ref{sect5} we analyze the different emission line components and study the physical properties of the different nebular regions probed by them. In Sects$.$ \\ref{incli} and \\ref{str+kin} we derive the mean nebular inclination and describe the kinematics of \\crl, respectively. The spatial distribution along the nebular axis of the extinction and the electron density are presented, respectively, in Sects$.$ \\ref{sext} and \\ref{physical}. In the latter, we also estimate the atomic and ionized mass in different regions. Finally, we discuss our results and give a possible scenario to explain the formation and future evolution of \\crl\\ in Sect$.$ \\ref{discuss}. The main conclusions of this work are summarized in Sect$.$ \\ref{conc}. ",
        "conclusions": "\\label{conc}  We have obtained ground-based, long-slit spectra in the ranges [6233-6806]\\,\\AA\\ and [4558-5331]\\,\\AA\\ for three different (1\\arcsec-wide) slit positions in the PPN \\crl. Based on the analysis of these spectra and direct (H$\\alpha$ and continuum) images, we find the following results:  $-$ The Balmer line emission results from the superposition of unscattered and scattered components. The unscattered emission is locally produced in the shocked lobes and in a inner, compact region labeled A' (Fig$.$\\,\\ref{f1}). The scattered radiation originates in the central \\ion{H}{2} region (and also maybe in region A') and is reflected by the dust in the shocked lobes and in the unshocked AGB envelope. Most forbidden lines are entirely produced in the lobes but there is also forbidden line emission (e.g$.$ [\\ion{Fe}{3}] and [\\ion{S}{3}]) arising in the \\ion{H}{2} region and region A'.   $-$ We identify the Wolf-Rayet bump at $\\sim$\\,6540\\,\\AA\\ superimposed on the scattered continuum. This bump, and a number of high-excitation transitions such as, e.g., \\ion{He}{2}$\\lambda$4686\\,\\AA, originate in the dense ($n_{\\rm e}$\\,$\\sim$\\,10$^5$-10$^7$\\,\\cm3) inner \\ion{H}{2} region, and are seen after being reflected by the nebular dust.   $-$ The spectrum of region A' is similar to that of the inner \\ion{H}{2} region (dominated by high-excitation lines) but remarkably different from the spectrum of the shocked lobes (dominated by low-excitation lines), and is consistent with $n_{\\rm e}$\\,$\\sim$10$^4$-10$^5$\\,\\cm3 and full ionization in this region. Thus, region A' very likely represents the outermost parts of the central \\ion{H}{2} region, i.e$.$ it is gas ionized by the UV stellar radiation rather than by shocks.  $-$ The shock velocity derived from the profiles of line emission arising in the shocked lobes is $\\gsim$\\,200\\,\\kms.  However, the line ratios in the lobes are consistent with the predictions of bow-shock models with a lower shock velocity (less than about 90\\,\\kms, and probably $\\sim$\\,75\\,\\kms).  $-$ We have measured an increase (decrease) of the [\\ion{O}{3}]$\\lambda$5007\\AA\\ ([\\ion{O}{1}]$\\lambda$6300\\AA) line flux by a factor $\\sim$\\,2.5 (0.7) with respect to observations performed $\\sim$\\,10\\,yr ago. Such variations could be due to a recent increase of the shock velocity.  $-$ The mean nebular inclination (with respect to the plane of the sky) is $i$=24\\degr$\\pm$6\\degr\\ and, very likely, less than 39\\degr.  $-$ The material in the lobes of \\crl\\ is moving away from the star at velocities up to $\\sim$\\,200\\,\\kms. The projected velocity of the gas shows a non-linear radial variation, resulting in a kinematical age that varies from $\\sim$100\\,yr near the nebular center to $\\sim$\\,400\\,yr at the tips, without correcting projection effects. For a mean inclination of $i$=24\\degr, the kinematical age of the optical lobes is $\\lsim$\\,180\\,yr.      $-$ The inner \\ion{H}{2} region and region A' seem to be expanding at moderate velocity, $\\lsim$\\,30\\,\\kms, much slower than the shocked lobes.   $-$ The extinction along the nebula is found to decrease from the center (A$_V$\\,$>$\\,10\\,mag) to the lobe tips ($\\sim$\\,2\\,mag); beyond the bright optical lobes, where an H$\\alpha$ scattered halo is visible, the extinction increases again. This increase indicates a large density contrast between the dust inside the optical lobes and beyond them, i.e$.$ in the AGB CSE unaltered by the shocks. This result is consistent with the lobes of \\crl\\ being cavities in the AGB CSE excavated by post-AGB winds.   $-$ The electron density, with mean values $n_{\\rm e}$\\,$\\sim$\\,[5-8]$\\times$10$^3$ and $\\sim$\\,8$\\times$10$^3$-10$^4$\\,\\cm3, for the east and west lobe, respectively, shows no systematic variation along the lobes.  $-$ The ionization fraction in the lobes is X$\\sim$\\,0.03-0.1, derived by comparison of different line ratios with predictions from planar shocks models. The total number density in the lobes is then $\\sim$10$^5$-10$^6$\\,\\cm3.  $-$ The mass of atomic gas in the east (west) lobe is $>$1.3$\\times$10$^{-4}$\\,\\ms\\ ($>$7.0$\\times$10$^{-5}$\\,\\ms).  The masses of ionized gas in the lobes are smaller: $\\sim$6$\\times$10$^{-5}$\\,\\ms, for the east lobe, and $\\sim$4$\\times$10$^{-5}$\\,\\ms\\ for the western lobe. The ionized mass in region A' is comparable to the ionized mass in the lobes or even higher, up to $\\sim$10$^{-4}$\\,\\ms.   $-$ The limb brightening of the optical lobes and the lack of a systematic spatial gradient in $n_{\\rm e}$ suggest that most of the line emission arises in a thin shell of shocked material rather than in gas filling the interior of the lobes. The thin-shell geometry is expected from radiative shocks shaping and accelerating the AGB circumstellar envelope. The radiative nature of the shocks is inferred from the short cooling time of the immediate post-shock gas, $\\lsim$\\,0.04\\,yr (relative to the dynamical time).  $-$ The post-AGB wind is probably currently active, shocking and heating the AGB circumstellar material, since the time required for the gas in the shocked lobes (presently partially ionized) to cool down significantly below 10$^4$\\,K is $\\lsim$\\,2\\,yr, much shorter than the age of the nebula."
    },
    "0206/astro-ph0206493_arXiv.txt": {
        "abstract": "We have identified the third known accretion-powered millisecond pulsar, XTE J0929$-$314, with the {\\em Rossi X-Ray Timing Explorer}. The source is a faint, high--Galactic-latitude X-ray transient ($d\\gtrsim 5$~kpc) that was in outburst during 2002 April--June.  The 185~Hz (5.4~ms) pulsation had a fractional rms amplitude of 3--7\\% and was generally broad and sinusoidal, although occasionally double-peaked.  The hard X-ray pulses arrived up to 770~$\\mu$s earlier than the soft X-ray pulses.  The pulsar was spinning down at an average rate of $\\dot\\nu=(-9.2\\pm 0.4)\\times 10^{-14}$ Hz~s$^{-1}$; the spin-down torque may arise from magnetic coupling to the accretion disk, a magnetohydrodynamic wind, or gravitational radiation from the rapidly spinning pulsar.  The pulsations were modulated by a 43.6~min ultracompact binary orbit, yielding the smallest measured mass function ($2.7\\times 10^{-7} M_\\odot$) of any stellar binary.  The binary parameters imply a $\\simeq 0.01 M_\\odot$ white dwarf donor and a moderately high inclination.  We note that all three known accreting millisecond pulsars are X-ray transients in very close binaries with extremely low mass transfer rates.  This is an important clue to the physics governing whether or not persistent millisecond pulsations are detected in low-mass X-ray binaries. ",
        "introduction": "Accretion-powered millisecond pulsars, the presumed progenitors of millisecond radio pulsars, have proven surprisingly elusive for 20 years.  The first example, the X-ray transient SAX J1808.4$-$3658 ($P_{\\rm spin}=$401 Hz, $P_{\\rm orb}=$2 hr), was identified as a millisecond pulsar only four years ago \\cite[]{wij98b,chak98d}.  This only deepened the puzzle of why more examples are not known, since SAX~J1808.4$-$3658 is very similar to many of the $\\simeq50$ neutron stars in low-mass X-ray binaries (LMXBs) that are not known pulsars \\cite[]{pc99}.  Earlier this year, a second system was detected, the X-ray transient XTE J1751$-$305 ($P_{\\rm spin}=$ 435 Hz, $P_{\\rm orb}=$42.4 min; Markwardt et al. 2002).  We report here on the discovery of a third example, again with a very short binary period.  The faint X-ray transient XTE~J0929$-$314 ($l=260\\fdg1$, $b=14\\fdg2$) was discovered by the All Sky Monitor (ASM) on the {\\em Rossi X-Ray Timing Explorer (RXTE)}\\/ in 2002 April \\cite[]{remillard02a}.  A brief scanning observation with {\\em RXTE}\\/ detected persistent 185~Hz pulsations \\cite[]{remillard02b}, and further timing revealed a circular, 43.6-min binary orbit \\cite[]{gmrc02b}.  The high Galactic latitude makes this source ideal for multiwavelength study.  Variable optical \\cite[]{gre02,cacella02} and radio \\cite[]{rupen02} counterparts were detected at the X-ray source position measured with the {\\em Chandra X-Ray Observatory} (Juett et al.  2002, in preparation), and \\ion{C}{3}/\\ion{N}{3} $\\lambda\\lambda$4640--4650 and H$\\alpha$ $\\lambda$6563 emission lines were reported in the optical spectrum \\cite[]{cs02}.  In this Letter, we present a detailed analysis of the {\\em RXTE}\\/ observations.  \\vspace*{0.2in} ",
        "conclusions": "From the presence of persistent millisecond pulsations over a wide luminosity range in \\src, we can infer an upper limit on the pulsar's surface dipole magnetic field strength of $B\\lesssim 1\\times 10^9\\,d_{\\rm 10kpc}$~G from accretion torque theory (Psaltis \\& Chakrabarty 1999).  This is consistent with the expectation that the neutron star is a recycled pulsar whose magnetic field has decayed during prolonged mass transfer (Bhattacharya \\& van den Heuvel 1991). It is interesting to note that the system has binary parameters that are extremely similar to those of the other recently discovered millisecond X-ray pulsar XTE J1751$-$305 (Markwardt et al. 2002) as well as the slow (7.6~s) accreting X-ray pulsar 4U 1626$-$67 (Middleditch et al. 1981; Schulz et al. 2001), pointing to a similar evolutionary path.  A puzzling aspect is that, unlike the two millisecond pulsars, the LMXB 4U~1626$-$67 has a strong ($3\\times 10^{12}$ G) magnetic field in addition to its slow spin period, which may indicate that its neutron star was formed recently through accretion-induced collapse (see Yungelson, Nelemans, \\& van den Heuvel 2002 and references therein).  We can estimate a crude lower bound on the distance to XTE~J0929$-$314 by first noting that the time-averaged mass transfer rate driven by gravitational radiation in this \\centerline{\\epsfxsize=8.5cm\\epsfbox{f3.eps}} \\figcaption{Pulse timing residuals in \\src. {\\it Top panel:} Fit residuals with and without the Keplerian orbit included. {\\it Bottom panel:} Fit residuals for a constant $\\nu$ plus Keplerian orbit. The best fit parabola to these residuals is also shown, indicating the need for a $\\dot{\\nu}$ term in the model.  The especially large deviations near MJD 52425 may be due to pulse shape variations.  Error bars represent $1\\sigma$ uncertainties. \\label{resid} } \\bigskip \\noindent 43.6~min binary is (Faulkner 1971) \\begin{equation} \\dot M_{\\rm GW} = 5.5\\times 10^{-12} \\left(\\frac{M_{\\rm x}}{1.4\\,M_\\odot}\\right)^{2/3} \\left(\\frac{M_{\\rm c}}{0.01\\,M_\\odot}\\right)^2 \\,M_\\odot\\mbox{\\rm\\,yr}^{-1} , \\end{equation} where $M_{\\rm x}$ and $M_{\\rm c}$ are the neutron star and companion masses, respectively.  If we assume that the mass accretion rate during the outburst is at least as large as $\\dot M_{\\rm GW}$, then our non-detection at the end of the outburst requires a distance $d\\gtrsim 6$~kpc. Also, given an outburst recurrence time of $\\gtrsim 6.5$~yr (the time since {\\em RXTE}\\/ was launched), the 2002 outburst fluence implies a mean mass transfer rate equal to $\\dot M_{\\rm GW}$ for $d\\gtrsim 5$~kpc.  We thus conclude that the source lies at least 1.2~kpc above the Galactic plane.  Our limits indicate that the mass accretion rate $\\dot M$ during the outburst was not more than a few percent of the Eddington critical rate.  The detection of spin-down during the outburst may provide an opportunity to test accretion torque theory for X-ray pulsars in the low-$B$ regime. The torque on a magnetic star from a prograde accretion disk tends to spin-up the star for sufficiently high $\\dot M$, while accretion will be centrifugally inhibited (the so-called ``propeller'' regime) for sufficiently low $\\dot M$. However, for an intermediate range of $\\dot M$, the positive material torque due to accretion may be dominated by a spin-down torque due to one of a variety of mechanisms, even while accretion persists. Two possible spin-down mechanisms include magnetic coupling of the accretion disk and the magnetosphere (Ghosh \\& Lamb 1979, 1991) and expulsion of a centrifugally driven magnetohydrodynamic wind (Anzer \\& B\\\"orner 1980; Arons et al. 1984; Lovelace et al. 1995). In either case, the torque should not exceed the characteristic accretion torque $\\dot M\\sqrt{GM_{\\rm x}r_{\\rm co}}$ (where $r_{\\rm co}$=50~km is the pulsar's corotation radius), which is consistent with our measured $\\dot\\nu$ for $d\\gtrsim 6$~kpc.  Another possible mechanism is gravitational radiation by the rapidly spinning pulsar (Bildsten 1998; Andersson, Kokkotas, \\& Stergioulas 1999; Levin 1999).  Some of these issues can be explored by examining the relation between $\\dot M$ and $\\dot\\nu$, which has been done extensively for high-$B$ neutron stars (e.g., Bildsten et al. 1997) but never for the low-$B$ case.  A detailed analysis of the $\\dot M$-$\\dot\\nu$ correlation will first require a better understanding of the pulse shape variations, in order to limit systematic uncertainties in the frequency history.  Such work is currently in progress.  White dwarfs and neutron stars accreting from a hydrogen-rich companion cannot evolve to binary periods below about 80~min, corresponding to the so-called ``period minimum'' observed for most cataclysmic variables and LMXBs (Paczynski \\& Sienkiewicz 1981; Rappaport, Joss, \\& Webbink 1982).  Ultracompact ($P_{\\rm orb}\\lesssim 80$~min) binaries like \\src\\ must therefore have a low-mass, hydrogen-depleted (and probably degenerate) donor (Nelson, Rappaport, \\& Joss 1986).  Our measured orbital parameters further constrain the nature of the donor in this system.   The pulsar mass function, $f_{\\rm X} =  2.7\\times10^{-7}\\ M_\\odot$, is the smallest presently known for any stellar binary. This mass function gives a minimum companion mass ($i=90^\\circ$) of $M_{\\rm c}=0.008$ $M_\\odot$ for $M_{\\rm x}=1.4\\ M_\\odot$; it also implies $M_{\\rm c}<0.03 M_\\odot$ (95\\% confidence) for a uniform a priori distribution in $\\cos i$ ($M_{\\rm x}=2 M_\\odot$).  The mass-radius relation for a Roche-lobe--filling donor in a 43.6~min binary is $R_{\\rm c} = 0.04 (M_{\\rm c}/0.01\\, M_\\odot)^{1/3}\\,R_\\odot$ \\cite[]{ffw72}.  As expected, this has no intersection with the theoretical mass-radius relation for very low-mass hydrogen-rich stars (i.e. brown dwarfs, Chabrier et al. 2000; cf. Bildsten \\& Chakrabarty 2001). A very low-mass cold helium white dwarf (Zapolsky \\& Salpeter 1969; Nelemans et al. 2001) is easily consistent for masses of $\\simeq 0.01\\,M_\\odot$, implying that the system has a relatively high (and thus probable) inclination. We note that a cold carbon white dwarf is not consistent with our measured orbital parameters (Lai, Abrahams, \\& Shapiro 1991).  It may be possible for a helium dwarf donor to retain a small residual hydrogen content (Podsiadlowski, Rappaport, \\& Pfahl 2002); this is of particular interest, given the report of an H$\\alpha$ $\\lambda$6563 emission line in the optical spectrum (Castro-Tirado et al. 2002).  Binary evolution theory predicts that many (if not most) of the $\\simeq$50 known neutron stars in LMXBs are spinning at millisecond periods (Bhattacharya \\& van den Heuvel 1991).  Of these, only three are now known pulsars, with persistent millisecond pulsations during their X-ray active states.\\footnote{ Another ten NS/LMXBs show millisecond oscillations during thermonuclear X-ray bursts; these are probably also signatures of rotation (see, e.g., Strohmayer \\& Markwardt 2002).} All three are low-luminosity transients in very close binaries, with extremely small time-averaged $\\dot M$; evidently, millisecond pulsations are easier to detect in such systems.   This supports the suggestion that magnetic screening by freshly accreted material may prevent the formation of persistent X-ray pulses in NS/LMXBs with $\\dot M$ above a critical value (Cumming, Zweibel, \\& Bildsten 2001)."
    },
    "0206/astro-ph0206170_arXiv.txt": {
        "abstract": "Multiple reflection of X-rays may be important when an accretion disc and its hot corona have a complicated geometry, or if returning radiation due to gravitational light bending is important, or in emission from a funnel such as proposed in some gamma-ray burst models. We simulate the effects of multiple reflection by modifying the boundary condition for an X-ray illuminated slab.  Multiple reflection makes the soft X-ray spectrum steeper (softer) and strengthens broad emission and absorption features, especially the K-shell features of iron.  This may be important in explaining the spectra of sources such as the Narrow-Line Seyfert 1 galaxy \\h. ",
        "introduction": "The hard X-ray emission from luminous Seyfert 1 galaxies is often modelled as due to an optically-thin corona above a dense accretion disc. This model successfully accounts for the reflection component often found in the spectra of these objects (Pounds \\etal\\ 1990). The dense optically-thick accretion disc is often assumed to be relatively cold, with iron only weakly ionized (e.g. George \\& Fabian 1991; Matt, Perola \\& Piro 1991; Magdziarz \\& Zdziarski 1995). Sometimes however it may be highly ionized by the irradiating X-rays, resulting in a more complex spectrum (Ross \\& Fabian 1993; \\.{Z}ycki \\etal\\ 1994; Ross, Fabian \\& Young 1999; Nayakshin, Kazanas \\& Kallman 2000; Ballantyne, Ross \\& Fabian 2001; R\\'{o}\\.{z}a\\'{n}ska \\etal\\ 2002). Generally the disc is assumed to be flat.  Recent numerical simulations of luminous accretion discs (Turner, Stone \\& Sano 2002), particularly when radiation pressure is important, suggest that they are clumpy, irregular and possibly corrugated. In this case, if the corona lies close to the disc, the solid angle subtended by the disc at the corona may be much higher than $2\\pi$, the value for a flat disc, and the reflection spectrum itself may be reflected. This will certainly happen if the disc is corrugated and the corona lies close to the bottom of corrugations. The observer may then not see any direct emission from the corona but only multiply-scattered flux.  (Rapid variability of an AGN requires that any such corrugations be on a relatively small spatial scale.)  We attempt here to simulate multiply-scattered reflection spectra for ionized discs in a simple generic manner using the code described by Ross \\& Fabian (1993). Multiply-scattered reflection spectra have been invoked as one explanation (Fabian \\etal\\ 2002) for the remarkable spectrum of the Narrow Line Seyfert~1 galaxy \\h\\ observed by XMM-Newton (Boller \\etal\\ 2002). They might also be relevant if the accretion disc extends well within 6 Schwarzschild radii so that returning radiation due to light bending by the strong gravity of the black hole is important (Cunningham 1975; Martocchia, Matt \\& Karas 2002). Multiple reflection is expected naturally if the accretion flow consists of clouds rather than a disc (e.g. Celotti, Fabian \\& Rees 1992; Collin-Souffrin \\etal\\ 1996; Malzac 2001). Multiple reflection could also be important in hypernova models for Gamma-ray bursts (M\\'{e}sz\\'{a}ros 2001), in which the emission is surrounded by ejecta.  We expect that multiple reflection will enhance both broad absorption and emission features in the spectrum, relative to single reflection spectra. The net effect is to produce a considerably steeper (softer) spectrum with pronounced broad spectral features around the K-shell emission of iron and other elements. ",
        "conclusions": "Multiple X-ray reflection creates a steep soft X-ray spectrum and a strong, broad, iron absorption/emission feature. The power lost from the harder energies above $\\sim 1\\ {\\rm keV}$ emerges at lower energies, principally in the 0.05--1~keV band. Overall the spectra resemble those from slightly higher ionization parameters and much higher abundances (Ballantyne, Fabian \\& Ross 2002).  A relativistically blurred version of the spectrum for $\\varepsilon=0.25$ and $\\xi=2000$ is shown in Figure~5.  That has been used to simulate what would be observed with the XMM-Newton EPIC pn CCD, and this is also displayed in Fig.~5 in the form of a ratio to an incident power-law of photon index 2.7. The result compares well to the spectrum of \\h\\ found by Boller \\etal\\ (2002). If we fit the simulated spectrum from 2--10~keV with a power-law plus a broad gaussian line, then we find that the best-fitting line has a central energy of 4.8~keV, a width of 1.5~keV and an equivalent width of about 6~keV. Positive residuals remain over the 6--7 keV band. It is not our intention to carry out a detailed fit to the data of \\h, which would require a comprehensive grid of models, but merely to confirm that multiple reflection may be relevant there, as suggested by Fabian et al. (2002). Since the spectrum of \\h\\ is unusual, multiple reflection may occur only occasionally among Narrow Line Seyfert 1 galaxies, and/or the primary continuum may be seen in addition to the reprocessed component most of the time.  Multiple reflection can be important when light bending is strong enough to cause part of the reflected spectrum to return to the disk. This occurs close to the centre of disks around maximally spinning Kerr black holes (Martocchia \\& Matt 1996; Dabrowski et al 1997; Dabrowski \\& Lasenby 2001; Martocchia, Matt \\& Karas 2002). Strong X-ray reflection has also been invoked (Vietri et al 2001; Ballantyne \\& Ramirez-Ruiz 2001) for the production of X-ray lines seen in some gamma-ray bursts (e.g. Piro et al 2000). If the emission is from a funnel geometry (Rees \\& Meszaros 2000; McLaughlin et al 2002) then the calculations presented here could be relevant. Lower abundances are implied than obtained from single reflection calculations."
    },
    "0206/astro-ph0206420_arXiv.txt": {
        "abstract": "The high resolution ($R$=25,000) infrared M band spectrum of the massive protostar \\irsnine\\ shows a narrow absorption feature at 4.779 \\mum\\ (2092.3 \\waven) which we attribute to the vibrational stretching mode of the \\thirteenco\\ isotope in pure CO icy grain mantles.  This is the first detection of \\thirteenco\\ in icy grain mantles in the interstellar medium.  The \\thirteenco\\ band is a factor of 2.3 narrower than the apolar component of the \\twelveco\\ band. With this in mind, we discuss the mechanisms that broaden solid state absorption bands. It is shown that ellipsoidally shaped pure CO grains fit the bands of both isotopes at the same time.  Slightly worse, but still reasonable fits are also obtained by CO embedded in N$_2$--rich ices and thermally processed O$_2$--rich ices. In addition, we report new insights into the the nature and evolution of interstellar CO ices by comparing the very high resolution multi-component solid \\twelveco\\ spectrum of \\irsnine\\ with that of the previously studied low mass source L1489 IRS. The narrow absorption of apolar CO ices is present in both spectra, but much stronger in \\irsnine. It is superposed on a smooth broad absorption feature well fitted by a combination of CO$_2$ and H$_2$O--rich laboratory CO ices. The abundances of the latter two ices, scaled to the total H$_2$O ice column, are the same in both sources. We thus suggest that thermal processing manifests itself as evaporation of apolar ices only, and not the formation of CO$_2$ or polar ices.  Finally, the decomposition of the \\twelveco\\ band is used to derive the \\twelveco/\\thirteenco\\ abundance ratio in apolar ices. A ratio of \\twelveco/\\thirteenco=71$\\pm$15 (3$\\sigma$) is deduced, in good agreement with gas phase CO studies ($\\sim$77) and the solid \\twelveco$_2$/\\thirteenco$_2$ ratio of 80$\\pm$11 found in the same line of sight. The implications for the chemical path along which CO$_2$ is formed are discussed. ",
        "introduction": "~\\label{se13co:intro}  Ever since the detection of interstellar solid CO \\citep{soif79, lacy84} in the 4.67 \\mum\\ spectra of protostars and background objects, the absorption band profile has been used as a diagnostic of the composition and evolution of interstellar ices \\citep{whit85, sand88, tiel91, chia98, teix98}. The absorption band was found to consist of a narrow feature accompanied by a broader feature at longer wavelengths. With the help of laboratory simulations it was found that the broad component is due to CO mixed with H$_2$O (`polar' ices) and the narrow feature due to pure CO or CO mixed with apolar species such as O$_2$, N$_2$ or CO$_2$. The relative depth of the apolar and polar ices is thought to reflect thermal processing in the envelopes of protostars, because these ices have quite different sublimation temperatures (18 K versus 90 K respectively). Also, the knowledge gained from observing the interstellar CO band is invaluable in studying the outgassing behavior of cometary ices as comets approach the sun \\citep{saal88}.  We have therefore started a program to measure the interstellar CO ice band at very high spectral resolution ($R=$25,000), more than an order of magnitude higher than customary until now, using the NIRSPEC spectrometer at the Keck II telescope. In Paper I of this program, we presented the spectrum of the low mass protostar L1489 IRS in the Taurus molecular cloud (Boogert, Hogerheijde, \\& Blake 2002).  At the high spectral resolution we discovered a new, third, CO component on the short wavelength side of the absorption band, which is compatible with absorption by CO$_2$--rich CO ices.  Combining the ice observations with information obtained from the gas phase CO lines in the same spectrum, we concluded that the CO ices are thermally processed in the upper layers of the circumstellar disk surrounding L1489 IRS.  In this Paper, we present the high resolution $M$ band spectrum of the massive protostar \\irsnine. This source has a rich and well studied infrared ice band absorption spectrum \\citep{whit96}. The ices are thought to reside in a thick and young circumstellar envelope, surrounding a modest hot core where the ices have evaporated \\citep{mitc90}.  Indeed, large gas phase depletion factors are needed in envelope models explaining millimeter wave emission lines toward \\irsnine\\ \\citep{tak00}. The contribution of ice absorption from unrelated cold foreground clouds in the NGC 7538 complex is likely small given the weakness of ice bands toward nearby protostars with more evolved envelopes (e.g. NGC 7538 : IRS1; \\citealt{lacy84}).  The CO ice band toward \\irsnine\\ has a large apolar component, tracing unprocessed ices in the cold envelope \\citep{sand88, tiel91, chia98}, and thus forms an interesting contrast with L1489 IRS. In addition to the \\twelveco\\ band we report the first detection of solid \\thirteenco\\ in \\irsnine\\ (and in the interstellar medium in general), providing an independent tracer of the composition of apolar ices. It also offers a reliable way of measuring the interstellar carbon isotope ratio, in follow up to measurements of the solid \\twelveco$_2$/\\thirteenco$_2$ ratio obtained with the {\\it Infrared Space Observatory} (ISO; \\citealt{boog00}), and to gas phase measurements (see \\citealt{wils94}).  This Paper is structured as follows. The observations and data reduction procedure are described in \\S 2. The 2092 \\waven\\ absorption band is identified with solid \\thirteenco\\ in \\S 3.1.1 using laboratory spectra. In \\S 3.1.2 we compare the apolar component of the \\twelveco\\ band with the \\thirteenco\\ band, explaining the factor 2.3 wider \\twelveco\\ band. After isolating the apolar component of the \\twelveco\\ band from the underlying broader absorption, we derive the interstellar solid \\twelveco/\\thirteenco\\ abundance ratio in \\S 3.2. The astrophysical implications of these results are discussed in \\S 4. A refined picture for the processing of CO ices is presented in \\S 4.1.  By comparing solid \\twelveco/\\thirteenco\\ and $^{12}$CO$_2$/$^{13}$CO$_2$ isotope ratios we discuss the chemical pathway leading to the formation of interstellar CO$_2$ in \\S 4.2. ",
        "conclusions": ""
    },
    "0206/astro-ph0206085_arXiv.txt": {
        "abstract": "{ We employ classical statistical methods of multivariate classification for the exploitation of the stellar content of the Hamburg/ESO objective prism survey (HES). In a simulation study we investigate the precision of a three-dimensional classification ($T_{\\mbox{\\scriptsize eff}}$, $\\log g$, [Fe/H]) achievable in the HES for stars in the effective temperature range $5200\\,\\mbox{K}<T_{\\mbox{\\scriptsize eff}}<6800\\,\\mbox{K}$, using Bayes classification. The accuracy in temperature determination is better than 400\\,K for HES spectra with $S/N>10$ (typically corresponding to $B_J<16.5$).  The accuracies in $\\log g$ and [Fe/H] are better than $0.68$\\,dex in the same $S/N$ range. These precisions allow for a very efficient selection of metal-poor stars in the HES. We present a minimum cost rule for compilation of complete samples of objects of a given class, and a rejection rule for identification of corrupted or peculiar spectra. The algorithms we present are being used for the identification of other interesting objects in the HES data base as well, and they are applicable to other existing and future large data sets, such as those to be compiled by the DIVA and GAIA missions. ",
        "introduction": "Ever since powerful computers and digital spectra have become available, there have been efforts to develop algorithms for automatic spectral classification \\citep[for a review on the early works see][]{Kurtz:1984}. The advantages of automated procedures as compared to manual classification are obvious. First of all, only a few experts are able to perform accurate manual classifications, and it was therefore sought to ``freeze'' this expert knowledge into computer programs. Such programs would allow to obtain \\emph{objective} classifications by \\emph{quantitative} criteria, and much larger data sets could be processed than by manual classification. The latter issue has become ever more demanding, with upcoming survey missions like DIVA\\footnote{\\texttt{http://www.ari.uni-heidelberg.de/diva/}}, NGST\\footnote{\\texttt{http://ngst.gsfc.nasa.gov/}}, or GAIA\\footnote{\\texttt{http://astro.estec.esa.nl/GAIA/}}. With all these satellites, it is planned to detect millions of objects, or even one billion objects in the case of GAIA.  {\\small \\begin{table*}[htbp] \\caption{\\label{performance_comparison} Comparison of automatic spectral classification performances. } \\begin{center} \\begin{tabular}{lllllllll}\\hline\\hline Method & Type of spectra & $\\lambda$ range & Disp. & $S/N$ & Types & $\\sigma_{\\mbox{\\scriptsize type}}$ & $\\sigma_{\\mbox{\\scriptsize LC}}$ & Reference\\rule{0ex}{2.5ex}\\\\\\hline PCA     & Slit/photoelectric & 3500--4000\\,{\\AA} & 10\\,{\\AA}/px & & A0--G0 & 1.16 & 0.85 & W83\\rule{0ex}{2.5ex}\\\\ Metric dist. & Slit/CCD & 3800--5190\\,{\\AA} & 67\\,{\\AA}/mm & & F8--G8 & 0.4 & & LS94\\\\ Metric dist. & Slit/CCD & 3500--5100\\,{\\AA} & 1--2\\,{\\AA}/px & & B0--F5 & 1.5 & & P94\\\\ ANN     & IUE & 1150--3200\\,{\\AA} & 2\\,{\\AA}/px & & O3--G5 & 1.11 & & VP95\\\\ Metric dist. & IUE & 1150--3200\\,{\\AA} & 2\\,{\\AA}/px & & O3--G5 & 1.38 & & VP95\\\\ ANN & Slit/Reticon & 5750--8950\\,{\\AA} & 7\\,{\\AA}/px & & A0--A9 & $0.42^{\\ast}$ & $0.15^{\\ast}$ & WTD95\\\\ ANN & Slit/Reticon & 5750--8950\\,{\\AA} & 7\\,{\\AA}/px & & O4--M6 & $1.26^{\\ast}$ & $0.38^{\\ast}$ & WTD97\\\\ ANN+PCA & Slit/CCD & 3510--6800\\,{\\AA} & 5\\,{\\AA}/px & & O--M & 2.34 &      & SGG98\\\\ Manual  & Objective prism, widened & 3800--5190\\,{\\AA} & 108\\,{\\AA}/mm & $>100$? & B2--M7 & $0.6^{\\ast\\ast}$ & $0.25^{\\ast\\ast}$ & H75--88\\\\ Metric dist. & Digitized objective prism & 3800--5190\\,{\\AA} & 1--3\\,{\\AA}/px & $>100$? & B & 1.14 & & LS94\\\\ ANN   & Digitized objective prism & 3800--5190\\,{\\AA} & 1--3\\,{\\AA}/px & $>100$? & B2--M7 & $0.82^{\\ast\\ast\\ast}$ & & BJIvH98\\\\ ANN   & Slit/CCD                  & 3850--4450\\,{\\AA} & 0.65\\,{\\AA}/px & $>20$ & F5--K5 & $0.57$--$0.64$ & & Setal01\\\\ Bayes & Digitized objective prism & 3200--5300\\,{\\AA} & 7--18\\,{\\AA}/px & 10--30 & F2-K0 & $<1.6$ & $<0.55$ & This work\\\\\\hline\\hline\\\\[-1.5ex] \\multicolumn{9}{l}{References: W83=\\cite{Whitney:1983}; LS94=\\cite{LaSala:1994}; P94=\\cite{Penprase:1994}; VP95=\\cite{Viera/Ponz:1995};}\\\\ \\multicolumn{9}{l}{\\hspace{4ex}WTD95=\\cite{Weaver/Torres-Dodgen:1995}; WTD97=\\cite{Weaver/Torres-Dodgen:1997} SGG98=\\cite{Singhetal:1998};}\\\\ \\multicolumn{9}{l}{\\hspace{4ex}Houk75--88=\\cite{Houk:1975}, \\cite{Houk:1978}, \\cite{Houk:1982}, \\cite{Houk/Smith-Moore:1988}; BJIvH98=\\cite{Bailer-Jonesetal:1998a};}\\\\ \\multicolumn{9}{l}{\\hspace{4ex}Setal01=\\cite{Snideretal:2001}}\\\\ \\multicolumn{9}{l}{$^{\\ast}$ Mean absolute deviation}\\\\ \\multicolumn{9}{l}{$^{\\ast\\ast}$ According to \\cite{vonHippeletal:1994a}}\\\\ \\multicolumn{9}{l}{$^{\\ast\\ast\\ast}$ 68\\,\\% quantile} \\end{tabular} \\end{center} \\end{table*} }  In the last decade, much progress was made in the field of automatic spectral classification, and it was demonstrated that computers are actually capable of performing this task \\citep[for a recent, comprehensive review see][]{Bailer-Jones:2001}. Using Kurtz' metric distance approach \\citep{Kurtz:1984}, \\cite{LaSala:1994} automatically classified digitized objective prism spectra from Houk's plates, with good results ($\\sigma=1.14$ MK-types). \\cite{Penprase:1994} used a similar approach, and applied it to slit spectra with similar spectral resolution and a slightly larger wavelength coverage (see Tab. \\ref{performance_comparison} for a comparison of the data used, and results obtained). The spectral type accuracy he reached for B0--F5 stars was a bit worse than that of LaSala; i.e., $\\sigma=1.5$ MK-types. However, as we will see below, it is very difficult to compare the performance of classification algorithms based on the results published in the literature, because (a) rarely ever is the signal-to-noise ratio ($S/N$) of the data documented, and the achievable classification accuracy depends critically on $S/N$; (b) different wavelength ranges and spectral resolutions were used; and (c) the algorithms were applied to stars in differing ranges of spectral type.  The influence of the latter on the achievable classification accuracy is nicely demonstrated by comparing the results of \\cite{Weaver/Torres-Dodgen:1995} with those of \\cite{Weaver/Torres-Dodgen:1997}. In the former paper, the authors report on supervised automatic classification of stars of spectral type A0--A9 with a multi-layer artificial neural network (ANN) with one hidden layer, trained with a back-propagation algorithm. They reached a mean absolute deviation of 0.42 spectral types and 0.15 luminosity classes. In the second paper, the ANN was applied to stars in the range O4--M6, and the mean absolute deviations were only 1.26 spectral types and 0.38 luminosity classes.  The results of Weaver \\& Torres-Dodgen have also shown that spectral classification in the near infrared can be done with the same accuracy as in the ``classical'' MK spectral range, with spectra of much lower resolution.  The resolution used by Weaver \\& Torres-Dodgen was only 7\\,{\\AA} per pixel, and their spectral range 5750--8950\\,{\\AA}. Their results are comparable to that achieved by others at three times higher spectral resolution in the optical or UV.  To continue with our brief review, in recent years, ANNs have been successfully used for supervised automatic spectral classification by a couple of groups. All of them used multilayer back-propagation networks (MBPNs). \\cite{Viera/Ponz:1995} automatically classified spectra of O3--G5 stars obtained with the International Ultraviolet Explorer (IUE; dispersion 2\\,{\\AA} per pixel) with an MBPN. The 1\\,$\\sigma$ error was 1.11 spectral types. They found their ANN classification to be superior to a classification with a metric distance method ($\\sigma=1.38$ types). The data used by \\cite{Singhetal:1998} were optical (3500--6800\\,{\\AA}) slit spectra with a dispersion of 5\\,{\\AA} per pixel. They used Principal Component Analysis (PCA) to pre-process their spectra, and reduce the number of input nodes. They obtained an accuracy of 2.34 spectral types over the full MK range (O--M). \\cite{Bailer-Jonesetal:1998a} used again Houk's plate material, digitized with the APM plate scanner, yielding a wavelength range of 3800--5190\\,{\\AA} and a dispersion of 1--3\\,{\\AA} per pixel. Their best ANN configuration classified these spectra with an error distribution having a 68\\,\\% quantile of 0.82 types, and the luminosity classification was correct for 95\\,\\% of the test sample spectra.  Recently, \\cite{Snideretal:2001} used a MBPN for derivation of the stellar parameters $T_{\\mbox{\\scriptsize eff}}$, $\\log g$ and [Fe/H] from moderate resolution (0.65\\,{\\AA} per pixel) spectra. Although their aim is to assign \\emph{continuous} parameter values to each spectrum, while we as well as the above mentioned authors carried out \\emph{discrete} classifications, we include their work in our review because Snider et al.  applied their technique to metal-poor stars, which is also the object type we are mainly concerned with in this paper.  Snider et al. report classification accuracies of $\\sigma_{T_{\\mbox{\\tiny eff}}}=135$--$150$\\,K, $\\sigma_{\\log g}=0.25$--$0.30$\\,dex and $\\sigma_{\\mbox{\\tiny [Fe/H]}}=0.15$--$0.20$\\,dex. However, it appears from the upper panel of their Fig. 4 that subgiants and horizontal branch stars have been excluded from the sample of stars they studied. A rough graphical analysis of their Fig. 4 reveals that unlike in real samples of stars emerging e.g. from wide-angle spectroscopic surveys, which \\emph{do} contain subgiants and horizontal-branch stars, their sample can be classified in $\\log g$ with a similar precision by dividing it into two classes ``by hand'', that is, assigning $\\log g=2.5$ to all stars with $T_{\\mbox{\\scriptsize eff}}<5000$\\,K, and $\\log g=4.5$ to all stars with $T_{\\mbox{\\scriptsize eff}}>5000$\\,K. Furthermore, it is questionable that there is any feature present in their set of spectra which does allow for a gravity classification, since they used continuum divided spectra. The Balmer jump, which is a gravity indicator in cool stars, is therefore removed. In conclusion, while Snider et al. succeeded in using ANNs for automated classification in $T_{\\mbox{\\scriptsize eff}}$ and [Fe/H], it remains to be demonstrated with a realistic sample that rectified moderate-resolution spectra indeed contain the information needed for a useful gravity classification.    ANN techniques and ``classical'' statistical methods such as Bayes and minimum cost rule classifications often perform equally well, in terms of e.g. minimising the total number of misclassifications. In the present work, we employ statistical methods, because their mathematical properties are well-studied, and the formulation of classification rules in the framework of mathematical statistics makes them very transparent.  Before we go into details of the methods we developed (Sect. \\ref{Sect:Autoclass}), we give a brief overview of the Hamburg/ESO Survey (HES) in Sect. \\ref{Sect:HES}, for better readibility. In Sect. \\ref{Sect:Performance} we investigate the classification performance for stars in the effective temperature range $5200\\,\\mbox{K}<T_{\\mbox{\\scriptsize eff}}<6800\\,\\mbox{K}$ achievable in the HES, by a simulation study. We summarize our conclusions in Sect. \\ref{Sect:Conclusions}. ",
        "conclusions": "\\label{Sect:Conclusions}  We have demonstrated that automatic spectral classification of turnoff stars, using ``classical'' statistical methods, is feasible in the HES with high accuracy. Our results suggest that it might be possible to determine the metallicity distribution function (MDF) of the galactic halo directly from a large sample of HES spectra. The MDF is an important constraint for models of galactic chemical evolution \\citep[see, e.g.,][]{Ikuta/Arimoto:1999,Oey:2000}.  The described methods are currently being applied to the large HES data base of digital spectra, in order to select interesting stellar objects in an automated fashion, and fully exploit the large scientific potential present in our data base.  Our algorithms can easily be adapted for automatic classification of other large data sets, e.g. those to be compiled by the DIVA and GAIA missions."
    },
    "0206/astro-ph0206091_arXiv.txt": {
        "abstract": "We present a catalogue of the 322 damped Lyman alpha absorbers taken from the literature.  All damped Lyman alpha absorbers are included, with no selection on redshift or quasar magnitude.  Of these, 123 are candidates and await confirmation using high resolution spectroscopy. For all 322 objects we catalogue the radio properties of the background quasars, where known.  Around 60 quasars have radio flux densities above 0.1 Jy and approximately half of these have optical magnitudes brighter than $V = 18$.  This compilation should prove useful in several areas of extragalactic/cosmological research. ",
        "introduction": "High resolution spectroscopy of quasars reveals large numbers of absorption lines due to neutral atomic hydrogen.  In some cases, the neutral hydrogen column density can be very large ($N_{\\rm HI}\\gapp10^{20}$ \\scm)\\footnote{In the catalogue some column densities less than this are to be found. This is because in order to produce the most comprehensive list, we have also included any Lyman alpha absorbers which are designated as possibly damped in the literature (see Section 2).} and gives rise to a heavily damped absorption feature (\\eg \\pcite{wtsc86}). The aim of this paper is to present a catalogue of these damped Lyman alpha absorption systems (DLAs) and, where known, the radio flux densities of the background quasar. The catalogue has been compiled from the large number of studies reported in the literature.  The nature of DLAs is still open to debate. Interpretations include galactic disks \\cite{wtsc86}, low surface brightness galaxies \\cite{jbm99} or dwarf galaxies \\cite{mmv97} intersecting the sight-line towards a background quasar. Despite these possible differing morphologies, DLAs provide a powerful cosmological probe. The highest redshift ($ z> 3.5$) DLAs may account for a large fraction of the baryons at high redshift, suggesting they reveal gas prior to the bulk of the star formation history of the universe (\\pcite{pmsi01} and references therein). On the other hand, recent work \\cite{lyp+02} seems to indicate star formation rates which continue to increase with increasing redshift up to the highest galaxy redshifts observed in the Hubble Deep Field. The discovery of a DLA with truly primordial abundances would have a major impact on our understanding of the early chemical evolution of the universe, and a crucial reality check on the ever-elusive population III. This will also be important for studies of primordial deuterium abundances (see below), since deuterium is destroyed, and never created, by star formation and evolution.  High resolution spectroscopy can be used to study high chemical abundances over a large redshift range. In particular, the difficult ionisation corrections required to derive meaningful chemical abundances in Lyman-limit absorbers (where $\\log N_{\\rm HI} > 17.2$ \\scm, so that they are optically thick to Lyman continuum radiation, \\eg \\pcite{lan91}) can be avoided using DLAs since the observed hydrogen is probably all neutral \\cite{twl+89,lwt+91}. Additionally, at high neutral hydrogen column densities, species such as Zn{\\sc \\,ii} and Cr{\\sc \\,ii} may become detectable, which are important since depletion onto dust grains is thought to be negligible for the former, whereas the latter remains in the solid phase. This allows both the study of abundances and depletion patterns/dust reddening (see \\pcite{pbh90,pksh97}). Some further reasons why DLAs are of interest are: \\begin{enumerate} \\item Studies of the higher order hydrogen Lyman series in DLAs can be used to investigate the primordial deuterium abundance \\cite{wcip91}.  The advantage of using DLAs is that the deuterium column density can be somewhat larger than typical Lyman forest absorbers.  This may help to discriminate against H{\\sc \\,i} interlopers mimicking the deuterium line.  Two recent observational studies \\cite{ddm01,pb01} report such D/H measurements. \\item Radio observations of quasars with a sufficiently high radio flux density can provide information complementary to that of the DLA observations: 21cm H{\\sc \\,i} measurements reveal more detailed kinematic information since line saturation is less severe and provide a direct spin temperature of the cool component of the gas. Different radio and optical morphology of the background quasar also provides the opportunity of observing along slightly different sight-lines through the same absorption complex, with the potential of learning about the relative sizes of optical/radio emission regions and the cloud size of the absorbing gas. In those rare cases where the host quasar has a sufficiently strong millimetre flux and a foreground molecular cloud occults the quasar \\cite{wc94,wc96}, a wealth of detailed chemistry is revealed \\cite{gpb+97,cw99}. \\item Studies of high redshift dust in DLAs gives a handle on the chemical evolution and star formation rates at various cosmological epochs \\cite{pfh99} through the contribution of dust to quasar spectral energy distributions (\\eg \\pcite{kvg+96,bcm+00,cbm+00,ocb+01}). \\item Certain heavy element transitions provide cosmological probes of special interest. For example, species with ground and excited state transitions sufficiently close to each other in energy provide a unique means of measuring the cosmic microwave background temperature at high redshift \\cite{bjl73,mbc+86,spl00}. \\item Finally, recent detailed studies of the relative positions of heavy element atomic optical transitions and comparison with present day (laboratory) wavelengths, suggests that the fine-structure constant ($\\alpha\\equiv e^2/\\hbar c$) may have evolved with time. Inter-comparing atomic optical transitions with H{\\sc \\,i} 21cm and molecular millimetre transitions may yield an order of magnitude over the already highly sensitive optical results \\cite{cs95,dwbf98,mwf+00}. However, very few such constraints are available due to the paucity of quasar absorption systems where 2 of the 3 types of transition (optical atomic, H{\\sc \\,i} 21cm or molecular millimetre) exist. \\end{enumerate} It is this last point which is of interest to us: As well as providing a comprehensive list of these objects for use by the astronomical community in general (Section 2), this catalogue allows us to shortlist those DLAs most likely to exhibit radio absorption lines in order to further constrain the variation in fine-structure constant. In the final sections we present the DLAs occulting radio-loud quasars along with any radio absorption features published and outline our future plans regarding the sample. ",
        "conclusions": "As mentioned in the introduction, we have compiled this catalogue since the comparison between optical and radio absorption lines can provide a considerably more precise determination of $\\Delta\\alpha/\\alpha$: To a first approximation, the ratio of two optical transition frequencies used in the many-multiplet method \\cite{dfw98b,wfc+98} is $\\frac{\\omega_1}{\\omega_2} \\propto 1+0.1\\alpha^2$. However, the ratio of the hyperfine neutral hydrogen (21\\,cm) to an optical resonance transition frequency is directly proportional to $\\alpha^2$ \\ie about 10 times larger. Thus, a substantial improvement in the determination of any variation of $\\alpha$ could be made by obtaining further statistics from optical {\\it and} 21 cm lines in cosmological absorbers. The limit on the variation of $\\alpha$ can be obtained by the comparison of the \\HI~ 21 cm line with any other optical or radio line (Section 3.2). However, by using redshifted 21cm \\HI~ together with $\\alpha$-sensitive species such as iron, zinc, chromium and nickel \\cite{dfw98b}, frequently seen in DLAs, we simultaneously maximise sensitivity and take advantage of the different signs of the frequency shifts due to $\\alpha$ variation to help minimise systematic effects \\cite{mwf+01b,wmf+01}.    A new systematic effect which applies to tests for $\\Delta\\alpha/\\alpha$ involving a \\HI~ \\& optical comparison involves the possible different spatial characteristics of the radio and optical quasar emission.  Large differences can result in the radio and optical light probing slightly different lines-of-sight.  However, we note that there are examples where the radio and optical emission is known to coincide spatially, and those cases are clearly of particular interest (Section 3.1). In order to minimise the spatial segregation problem, the most reliable tests will come from comparing \\HI~ lines with neutral atomic or molecular species, or singly ionised species where the ionisation potential is smaller than that for neutral hydrogen.  \\subsection{Radio-loud quasars illuminating DLAs}  Of the known radio-loud ($S_{\\rm radio}\\gapp0.1$ Jy) systems, we summarise the current state of searches for atomic and molecular hydrogen (Section 3.2) absorption features. Note that with regard to the spatial distribution of the optical and radio emission, from the NVSS catalogue \\cite{ccg+98}, {\\it unless otherwise stated, the 1.4 GHz emission extends to a radius of $\\approx1'$ and the peak emission coincides with the given optical position (Table 1).}  \\noindent {\\bf PKS 0118--272:} A BL Lac object where \\scite{kc01a} failed to detect \\HI~ absorption at $z=0.5579$.  \\noindent {\\bf [HB89] 0149+336:} A gravitational lens candidate for which we could find no reference to radio absorption features.  \\noindent{\\bf PKS 0201+113:} A gravitational lens where \\scite{dob96,bbw97} have detected \\HI~ absorption at $z=3.388$.  \\noindent{\\bf [HB89] 0201+365:} No reference to radio absorption features found.  \\noindent{\\bf [HB89] 0215+015:} A BL Lac object where \\scite{bw83} failed to detect \\HI~ absorption.  \\noindent{\\bf [HB89] 0235+164:} A BL Lac object where \\HI~ absorption at $z=0.524$ has been detected (\\pcite{wbd82}; \\pcite{bw83}). \\scite{drrw92,wc95} failed to detect CO at the absorption redshift.  \\noindent{\\bf [HB89] 0248+430:} \\scite{lb01} have detected the \\HI~ absorption at the DLA redshift.  \\noindent {\\bf [HB89] 0329--255:} No reference to radio absorption features found.  \\noindent {\\bf  [HB89] 0335--122:} No \\HI~ absorption detected \\cite{kc02}.  \\noindent {\\bf PKS 0336--017:} \\scite{cld+96,kc02} failed to detect \\HI~ absorption at $z=3.0619$.  \\noindent {\\bf QSO 0347--211:} No reference to radio absorption features found.  \\noindent {\\bf PKS 0405--331:} As above.  \\noindent {\\bf QSO 0432--440:} As above. {\\it No NVSS data available.}  \\noindent {\\bf [HB89] 0438--436:} \\scite{dcw96} failed to detect CO in the torus of this AGN ($z=2.852$). {\\it No NVSS data available.}  \\noindent {\\bf [HB89] 0439--4319:} Tentative \\HI~ absorption detected by \\scite{kcsp01} in this low redshift source. {\\it No NVSS data available.}  \\noindent {\\bf PKS 0454+039:} No \\HI~ \\cite{bw83} nor H$_2$ \\cite{gb99} absorption has been detected. {\\it No optical/radio offset, but there is a second $30''$ radius radio source centered at 5 s to the West.}  \\noindent {\\bf [HB89] 0458--020:} In this blazar, \\scite{wbt+85,bwl+89} have detected \\HI~ absorption at $z=2.03945$. No H$_2$ nor CO (\\ie molecular) absorption has been detected \\cite{wc94a,gb99}.  \\noindent {\\bf PKS 0528--250:} \\scite{cld+96} failed to detect \\HI~ absorption at $z=2.8110$, although H$_2$ absorption in this DLA \\cite{fcb88,sp98,gb99} and CO emission in the $z=2.14$ DLA \\cite{bv93} have been detected. Note that no H$_2$ or CO absorption in either DLA was detected by \\scite{wc94a,lsb99}.  \\noindent {\\bf [HB89] 0537--286:} No reference to radio absorption features found.  \\noindent {\\bf [HB89] 0552+398:} Although Galactic H{\\sc \\,i} \\cite{dkvh83} and HCO$^+$ \\cite{ll96} absorption have been observed towards this quasar, no reference to absorption at the DLA (or any cosmological) redshift could be found.  \\noindent {\\bf J074110.6+311200:} In this optically variable quasar, \\scite{lsb+98,kgc01} have detected \\HI~ absorption at $z=0.2212$.  \\noindent {\\bf FBQS J083052.0+241059:} In this blazar, \\scite{kc01a} have detected \\HI~ absorption at $z=0.5247$.  \\noindent {\\bf IRAS F08279+5255:} A gravitational lens in which \\scite{cmo99} have detected CO 4$\\rr$3 emission at $z=3.911$, the redshift of the source. {\\it There is a weak central radio source at optical position with two stronger diagonally opposing sources near} 08h31m50s/52d$43'30''$ {\\it and} 08h31m25s/52d$46'30''$.  \\noindent {\\bf [HB89] 0834--201:} No reference to radio absorption features found for this blazar.  \\noindent {\\bf QSO 0913+003:} No reference to radio absorption features found.  \\noindent {\\bf QSO 0933--333:} As above. {\\it Offset from optical position at} 09h35m08.6s/-33d$32'34''$.   \\noindent {\\bf [HB89] 0938+119:} No reference to radio absorption features found. {\\it No offset but there is a second source to the South East near} 09h41m20.5s/11d$45''00'$.  \\noindent {\\bf [HB89] 0952+179:} \\scite{kc01a} have detected \\HI~ absorption at $z=0.2378$.  \\noindent {\\bf [HB89] 0957+561:} A gravitational lens where no \\HI~ absorption has been detected \\cite{kc02}.  \\noindent {\\bf [HB89] 1017+109:} No reference to radio absorption features found. {\\it Radio position offset $\\approx20''$ to the West of the optical centre.}  \\noindent {\\bf [HB89] 1021-006:} No reference to radio absorption features found for this optically variable quasar.  \\noindent {\\bf RX J1028.6-0844:} No \\HI~ absorption detected \\cite{kc02}.  \\noindent {\\bf PKS 1055--301:} No reference to radio absorption features found. {\\it Radio position offset $\\approx1'$ to the West of the optical centre.}  \\noindent {\\bf 2MASSi J1124428--170517 :}  No reference to radio absorption features found. {\\it Offset slightly from optical position at} 11h24m41.5s/-17d$05'10''$.  \\noindent {\\bf [HB89] 1127--145:} \\scite{lsb+98,ck00} have detected variable \\HI~ absorption at $z=0.3127$ towards this blazar.  \\noindent {\\bf [HB89] 1157+014:} \\scite{wbj81,bw83} have detected  \\HI~ absorption at $z=1.94362$.  \\noindent {\\bf LBQS 1213+0922:} No reference to radio absorption features found.  \\noindent {\\bf FBQS J121732.5+330538 :} \\scite{wc94a} failed to detect CO absorption at $z=1.9984$.  \\noindent {\\bf FBQS J122824.9+312837:} \\scite{bw83} failed to detect \\HI~ absorption at $z=1.7945$.  \\noindent {\\bf PKS B1228--113:} No reference to radio absorption features found.  \\noindent {\\bf QSO 1230--101:} As above.  \\noindent {\\bf PKS 1251--407:} As above. {\\it No NVSS data available.}  \\noindent {\\bf 3C 286:} \\HI~ at $z=0.69215$ by \\scite{br73} but no H$_2$ absorption has yet been detected \\cite{gb99}.  \\noindent {\\bf [HB89] 1331+170:} A blazar where \\scite{wd79,bw83} have detected \\HI~ absorption at $z=1.7764$, but \\scite{lsb99} failed to detect CO absorption.  \\noindent {\\bf PKS B1354--107:}. No \\HI~ absorption detected \\cite{kc02} in the $z_{abs}=2.966$ DLA. {\\it Radio position offset $\\approx15''$ to the West of the optical centre.}   \\noindent {\\bf [HB89] 1354+258:} No reference to radio absorption features found.  \\noindent {\\bf [HB89] 1402+044:} No reference to radio absorption features found for this BL Lac.  \\noindent {\\bf PKS B1418--064:} No reference to radio absorption features found.  \\noindent {\\bf [HB89] 1451--375:} \\scite{ck00} failed to detect \\HI~ absorption in this HST source.  \\noindent {\\bf 3C 336:} No reference to radio absorption features found for this optically variable quasar.  \\noindent {\\bf PMN J2130--4515:} No reference to radio absorption features found. {\\it No NVSS data available.}  \\noindent {\\bf PMN J2134--0419:} No reference to radio absorption features found.  \\noindent {\\bf [HB89] 2136+141:} No reference to radio absorption features found.  \\noindent {\\bf 3C 446:} A blazar not detected in \\HI~ absorption at the DLA \\cite{ck00} nor CO absorption at the quasar redshift \\cite{dcw96}.  \\noindent {\\bf QSO 2311--373:} No reference to radio absorption features found.  \\noindent {\\bf [HB89] 2314--409:} As above. {\\it No NVSS data available.}  \\noindent {\\bf MG3 J234456+3433:} \\scite{cld+96,kc02} failed to detect \\HI~ absorption at $z=2.9084.$  \\noindent Finally, note that \\HI~ absorption has been observed in the inferred (from metal lines) DLAs {\\bf 3 C196}, {\\bf LBQS 1229--0207} \\cite{wlfc95} and {\\bf [HB89] 1243--072} \\cite{lb01}.   \\subsection{Searching for new radio absorbers} If we summarise the current \\HI~ absorption results for the DLAs (Table 2), we see that although many of the positive results have very high column densities, this does not appear to be a prerequisite for \\HI~ absorption (\\ie FBQS J083052.0+241059). Perhaps also of relevance is the spectral energy distributions (SEDs): Note that all of the GHz peaked sources have high column densities and have all been detected in \\HI. Of the two inverted SEDs, one DLA has a high column density whereas the other is relatively low and \\begin{table} \\begin{center} \\begin{tabular}{l r c c c l} \\hline\\noalign{\\smallskip} Quasar & $\\tau$ & $\\log N_{\\rm HI}$  & $S$ & S.I. & Notes\\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} PKS 0118--272 &  $<0.007$ & 20.3 & 1.2 & 0.1 & \\\\ PKS 0201+113 & 0.09,0.04 & 21.3 & 0.3 & -- & GPS (2.6)\\\\ ${\\rm[HB89]}$  0215+015 & $<0.04$ & 19.9 & 0.9 &-- &  See caption \\\\ ${\\rm[HB89]}$  0235+164 & 0.05-0.5 & 21.6 & 1.8 & -0.2 & Inverted \\\\ ${\\rm[HB89]}$  0248+430 & 0.20 & 21.6 & 1.2 & -- & GPS (2.5)\\\\ ${\\rm[HB89]}$  0335--122 & $<0.008$ & 20.8 & 0.8& 0.3 & \\\\%straightforward [HB89] screws up table PKS 0336--017 &$<0.005$  & 21.2 & 1.3 & 0.6 & \\\\ ${\\rm[HB89]}$  0439--4319 & $<0.007$ & 20.0 & 0.4 & 0.2 & \\\\ PKS 0454+039 & $<0.01$& 20.7 & 0.4 & -- &  See caption \\\\ ${\\rm[HB89]}$  0458--020  & 0.3 & 21.7 & 2.5 & 0.3 & \\\\ PKS 0528--250 & $<0.2$ & 21.2 & 1.9 & 0.5 & For $z_{abs}=2.811$ DLA \\\\ J074110.6+311200 & 0.07 & 21.2 & 1.9 & -- & GPS (2.9)\\\\ & & & &  &$z_{abs}=0.221$ DLA \\\\ FBQS J083052.0+241059 & 0.007 & 20.3 & 0.8 & -0.2 & Inverted\\\\ ${\\rm[HB89]}$  0952+179 &0.013 & 21.3 & 1.2 & 0.3 & \\\\ ${\\rm[HB89]}$  0957+561 & $<0.004$ & 20.3 & 0.9 & 1.3 & \\\\ RX J1028.6-0844  &  $<0.03$ & 20.1 & 1.7 & 0.9 & $z_{abs}=3.42$ DLA \\\\ ${\\rm[HB89]}$  1127--145 & 0.06 & 21.7 &6.2 &-- & GPS (1.4)\\\\ ${\\rm[HB89]}$  1157+014 & 0.05 & 21.8 & 1.0 & 0.8 & \\\\ FBQS J122824.9+312837 & $<0.05$ & 19.1 & 0.3 & 0 & \\\\ 3C 286 & 0.11 & 21.3 & 19.0 & 0.6 & \\\\ ${\\rm[HB89]}$  1331+170 & 0.020 & 21.4 & 0.6 & -- &  See caption \\\\ PKS B1354--107 & $<0.05$ &20.8 & 0.2 & 0 & $z_{abs}=2.996$ DLA \\\\ ${\\rm[HB89]}$  1451--375 & $<0.006$ & 20.1 & 1.8 & 0.2 & \\\\ 3C 446 &$<0.02$  & 20.9 & 7.4 & 0.5 & \\\\ MG3 J234456+3433  &$<0.04$  & 21.2 & 0.3 & 0.2 & \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\label{sum} \\addtocounter{table}{1} % \\caption{The radio-loud DLAs in which \\HI~ absorption has been searched for.  $\\tau$ is the optical depth of the \\HI~ line, with $3\\sigma$ upper limits quoted, as given by the references in Section 3.1. For PKS 0201+113 the values are from \\protect\\scite{dob96} and \\protect\\scite{bbw97}, respectively. PKS 0336--017 and MG3 J234456+3433 these are \\protect\\scite{kc02} results; \\protect\\scite{cld+96} obtained $\\tau<0.02$ and 0.1, respectively. $S$ is the approximate flux density in Janskys at $z_{abs}$ and S.I. is the spectral index (both are estimated from the flux density values in Table 1 and $S\\propto\\nu^{-{\\rm S.I.}}$). In the last column, GPS designates a GHz peaked source with the approximate turnover frequency given in parenthesis.  In the case of the ``U-shaped'' SEDs: ${\\rm[HB89]}$ 0215+015 is known to exhibit radio outbursts (\\eg \\protect\\pcite{lo85}) and so the flux densities quoted will be variable. For PKS 0454+039 and ${\\rm[HB89]}$ 1331+170, these could be due to an anomalous flux density measurement and both are considered flat spectrum sources (\\eg \\protect\\cite{wgbb84,msm+97}).} \\end{center} \\end{table} both of the flat SED detections have high column densities. Finally, the two steep spectrum quasars which illuminate DLAs detected in \\HI~ absorption ([HB89] 1157+014 and 3C 286) also have high column densities.  Because of the relation between turnover frequency and source size \\cite{ffs+90,ob97}, we may expect a higher \\HI~ absorption detection rate from flat and inverted SED sources, since these result from similar optical and radio lines-of-sight.  However, as it stands, the statistics are too small (Table 2) and so in order to maximise our sample, it appears that the way to proceed is an unbiased search for H{\\sc\\,i} in the DLAs occulting the remaining radio-loud quasars.  As mentioned in Section 1, as well as optical and \\HI~ comparisons, the inter-comparison of atomic and molecular lines will also give a ten-fold increase in accuracy for $\\Delta\\alpha/\\alpha$: Due to its zero dipole moment and small moment of inertia, molecular hydrogen cannot be directly observed at radio frequencies\\footnote{In the case of $z>1.8$ sources, however, the ultra-violet lines of H$_2$ are redshifted into the optical window, making molecular hydrogen readily observable at these frequencies. As well as for PKS 0528--250 (Section 3.1) molecular hydrogen has also been detected in the DLAs occulting the radio-quiet quasars [HB89] 0000-263 \\cite{lmc+00a}, LBQS 0013--0029 \\cite{gb97,psl02}, [HB89] 0347--383 \\cite{lddm01}, LBQS 1232+0815 \\cite{gb97,spl00} and the inferred \\cite{wlfc95} DLA [HB89] 0551--366 \\cite{lsp02}.} and so it is the usual practice infer the presence of this from the millimetre rotational lines of such molecules as CO. In order to also take advantage of this, we have applied for time to search for molecular absorption lines in the DLAs occulting mm-loud quasars with the IRAM 30 metre and Swedish ESO Sub-millimetre telescopes. Recently (April 2002), we have been awarded time on the Australia Telescope Compact Array in order to obtain 90 GHz flux measurements for the whole radio-loud sample, as a means of selecting new sources in which to search for millimetre absorption. The results will be published in forthcoming papers.  Also, with regard to finding new systems in which there may be absorption (in all three frequency regimes), we see that 13 of the quasars are known to be BL Lac/optically variable/blazars\\footnote{BL Lacs and Optically Violent Variables are known collectively as blazars. In these radio-loud active galactic nuclei the radio jet is relativistically beamed close to the line-of-sight (\\eg \\pcite{pet97}).} and that 3 of the sources are known gravitational lenses.  This may be of interest as of the four known high redshift millimetre (\\ie molecular) absorbers, two are BL Lac objects; B 0218+357 \\cite{wc95} and PKS 1413+135 \\cite{wc94}\\footnote{The former, as well as PKS 1830--211, is also a gravitational lens \\cite{wc96}.}. This may suggest several strategies for finding similar new absorbers \\cite{sr99} which could prove useful in appending to this catalogue."
    },
    "0206/astro-ph0206058_arXiv.txt": {
        "abstract": "We have mapped the $^{12}$CO $J$=1-0 and $J$=2-1 line emission in Mrk~86, one of the most metal-deficient Blue Compact Dwarf galaxies so far detected in $^{12}$CO. The $^{12}$CO emission is distributed in a horseshoe-like structure that follows the locus of the most recent star formation regions. The minimum in molecular-line emission corresponds to the position of an older, massive nuclear starburst. The H$_{2}$ mass of the galaxy (in the range 0.4-5$\\times$10$^{7}$\\,M$_{\\odot}$) and its morphology have been compared with the predictions of hydrodynamic simulations of the evolution of the interstellar medium surrounding a nuclear starburst. These simulations suggest that the physical conditions in the gas swept out by the starburst could have led to the formation of the ring of molecular gas reported here. This result provides an attractive scenario for explaining the propagation (in a galactic scale) of the star formation in dwarf galaxies. ",
        "introduction": "\\label{sec3}  In order to study the physical conditions in the molecular gas we have derived the 2-1/1-0 brightness temperature ratio in Mrk~86.  The mean values of these ratios are 1.34$\\pm$0.46 and 0.40$\\pm$0.14 under {\\it Uniform Filling} and {\\it Point Source} approximations, respectively, using only those positions where both transitions are clearly detected. Although the {\\it Point Source} approximation is usually assumed for BCD galaxies (Sage et al$.$ 1992; Meier et al$.$ 2001), our CO maps suggest that the emission in Mrk~86 arises in an extended structure formed by several (most likely unresolved) clouds. Thus, we have decided to use a more precise approach for deriving the 2-1/1-0 brightness temperature ratio. We deconvolve our CO(1-0) map with a 22\\arcsec\\ gaussian beam. Then, we determine the 1-0 and 2-1 line intensities at each of the positions detected in both transitions along with the size of the emitting regions (in the deconvolved map) inside the 1-0 and 2-1 beams for each of these positions. The average of the 2-1/1-0 brightness temperature ratios derived at each point yields a 2-1/1-0 brightness temperature ratio of 1.06$\\pm$0.40. We have not found any systematic difference between the 2-1/1-0 ratio at the position of the nuclear starburst and the outermost regions of the galaxy, which indicates that there is no significant change of the excitation conditions.  Meier et al$.$ (2001) failed to detect $^{12}$CO(3-2) at the center of Mrk~86 and provided a very low upper limit of 0.45 to the 3-2/1-0 brigthness temperature ratio using the 1-0 intensity published by Sage et al$.$ (1992). However, the value reported by Sage et al$.$ at this position, 1.12\\,K\\,km\\,s$^{-1}$, is a factor of two larger than that measured by us, 0.56\\,K\\,km\\,s$^{-1}$. This difference could be due to calibration or pointing inaccuracies. Therefore, we have decided to adopt a more conservative upper limit of 0.9 to the 3-2/1-0 ratio.  Under the LTE and optically thick assumptions we can estimate the excitation temperature ($T_{\\mathrm{ex}}$) of the molecular gas from the 2-1/1-0 and 3-2/1-0 brightness temperature ratios. In particular, from 2-1/1-0 we infer $T_{\\mathrm{ex}}$$>$6\\,K and, according to the 3-2/1-0 upper limit, $T_{\\mathrm{ex}}$$<$45\\,K. From the predictions of Large Velocity Gradient models and the 2-1/1-0 and 3-2/1-0 ratios measured we have also derived a range of compatible molecular gas densities ($n_{\\mathrm{H}_{2}}$) and kinetic temperatures ($T_{\\mathrm{kin}}$). A [CO/H$_{2}$] abundance of 8$\\times$10$^{-6}$ ($Z_{\\odot}$/10) and a velocity gradient of 1\\,km\\,s$^{-1}$\\,pc$^{-1}$ are assumed (see e.g$.$ Meier et al$.$ 2001). The line ratios measured suffer from a strong degeneracy in the physical properties of the molecular gas, and only rough constraints can be obtained. We can only conclude that, if $n_{\\mathrm{H}_{2}}$$>$500\\,cm$^{-3}$, then the molecular gas cannot be warmer than $T_{\\mathrm{kin}}$=40\\,K.  We have also obtained a variety of estimates of the mass of molecular gas in Mrk~86. First, we have computed the molecular mass under the optically-thin approximation from the $^{12}$CO(1-0) integrated luminosity ($L_{\\mathrm{CO}}$ = 2$\\times$10$^{6}$\\,K\\,km\\,s$^{-1}$\\,pc$^{2}$) adopting an average value for the excitation temperature ($<$$T_{\\mathrm{ex}}$$>$ = 20\\,K). This yields $M_{\\mathrm{thin}}$ = 4$\\times$10$^{6}$\\,M$_{\\odot}$. It is worth noting that this transition is rarely optically thin ($^{13}$CO and C$^{18}$O transitions would be more appropriated) and so the value given above provides only a rough but firm lower limit to the actual molecular mass of the galaxy. From the $^{12}$CO(1-0) line width ($\\Delta\\,v_{1/2}$ = 75\\,km\\,s$^{-1}$) we can also estimate the molecular gas mass assuming virial equilibrium, which leads to $M_{\\mathrm{vir}}$ = 7$\\times$10$^{8}$\\,M$_{\\odot}$. Again, the value derived, although represents a firm upper limit, is very far from the actual molecular mass since the gas velocities measured are supported by rotation. Moreover, a large fraction of the mass in the galaxy central regions is in the form of stars (GZG99).  An estimate of the H$_{2}$ mass can be also obtained using a $L_{\\mathrm{CO}}$-to-$M_{\\mathrm{H}_2}$ conversion factor appropriate to the low metallicity of Mrk~86 ($X_{\\mathrm{CO}}$$\\equiv$$M_{\\mathrm{H}_2}$/$L_{\\mathrm{CO}}$; see e.g$.$ Barone et al$.$ 2000). The dependence of this factor on metallicity was studied by Arimoto et al$.$ (1996; A96) using the virial masses and CO luminosities of Giant Molecular Clouds in nearby galaxies. These authors proposed the following relation, log\\,($X_{\\mathrm{CO}}$/$X_{\\mathrm{MW}}$)=$-$[O/H]$+$8.93, where $X_{\\mathrm{MW}}$ is the Milky Way conversion factor, 1.56$\\times$10$^{20}$ molecules\\,cm$^{-2}$\\,(K\\,km\\,s$^{-1}$)$^{-1}$ (Hunter et al$.$ 1997). The average value that we adopt for the gas in Mrk~86 is 1/10$^{\\mathrm{th}}$ the Solar value, which yields $X_{\\mathrm{CO}}$ = 1.6$\\times$10$^{21}$ molecules\\,cm$^{-2}$\\,(K\\,km\\,s$^{-1}$)$^{-1}$. From this value and the CO luminosity given above we derive a molecular gas mass (helium excluded) of $M_{\\mathrm{H}_2}$=5$\\times$10$^{7}$\\,M$_{\\odot}$.  Based on recent studies of low metallicity BCD galaxies we believe, however, that the A96 relation is probably over-estimating the molecular gas mass in those low-metallicity BCDs detected in CO, and, in particular, the molecular mass given above. Observations of the SMC carried out by Rubio et al$.$ (1993) have shown that the $L_{\\mathrm{CO}}$-to-$M_{\\mathrm{H}_2}$ conversion factor depends on the spatial scale of the CO emitting region considered: at spatial scales of 10-20\\,pc $X_{\\mathrm{CO}}$ is only slightly higher than $X_{\\mathrm{MW}}$ while at larger scales ($\\sim$100\\,pc) its value increases dramatically. This suggests the presence of large amounts of H{\\sc i} or {\\it hidden} H$_{2}$ (not associated with CO) between the dense clumps where the CO emission arises. The lack of detection of diffuse H$_{2}$ in I~Zw~18 by FUSE (Vidal-Madjar et al$.$ 2000) and the very low upper limit derived to the [CII]/CO line ratio in I~Zw~36 (Mochizuki \\& Onaka 2001) suggest that, in the case of the low-metallicity BCDs, most of the H$_{2}$ is probably in the form of these dense clumps. In dwarf irregular galaxies like IC~10 the intense [CII] emission observed has been argued as due to the presence of large amounts of H$_{2}$ with no CO emission associated (Madden et al$.$ 1997). These results suggest that in low-metallicity BCDs most of the H$_{2}$ is probably associated with emitting CO and, therefore, the global $L_{\\mathrm{CO}}$-to-$M_{\\mathrm{H}_2}$ conversion factor would be similar to the local value derived from the analysis of the CO line ratios. Noteworthly, Large Velocity Gradient modeling of the few low-metallicity BCDs detected in $^{12}$CO favors local conversion factors lower than the A96 predictions and similar, in some cases, to the Galactic value (Sage et al$.$ 1992; Barone et al$.$ 2000). Unfortunately the lack of observations on optically-thin $^{13}$CO and C$^{18}$O lines, which would provide an accurate determination of the physical conditions in the gas, makes difficult to obtain definitive conclusions in this sense. Finally, it is also worth noting that other authors have proposed a shallower dependence of the $L_{\\mathrm{CO}}$-to-$M_{\\mathrm{H}_2}$ conversion factor on metallicity than that argued by Arimoto et al$.$ (1996). For example, using the relation proposed by Wilson (1995) we derive a molecular mass of $M_{\\mathrm{H}_2}$=2$\\times$10$^{7}$\\,M$_{\\odot}$.  Therefore, we conclude that the H$_{2}$ mass of Mrk~86 is certainly in the range 0.4-70$\\times$10$^{7}$\\,M$_{\\odot}$ (based on the {\\it optically-thin} and {\\it virial} approximation estimates) with a most probable value in the range 0.4-5$\\times$10$^{7}$\\,M$_{\\odot}$. The latter range is given by the uncertainties expected in using the metallicity-scaled $L_{\\mathrm{CO}}$-to-$M_{\\mathrm{H}_2}$ conversion factor. ",
        "conclusions": ""
    },
    "0206/astro-ph0206328_arXiv.txt": {
        "abstract": "Strong evidence for deep mixing has been uncovered for slowly rotating F, and A stars of the main sequence. As the accretion/diffusion model for the formation of \\lboo\\ stars is heavily dependent on mixing in superficial regions, such deep mixing may have important repercussions on our understanding of these stars. It is shown that deep mixing at a level similar to that of FmAm stars increases the amount of matter that needs to be accreted by the stars with respect with the standard models by some three orders of magnitude. It is also shown that significantly larger accretion rates have to be maintained, as high as $10^{-11}$~M$_\\sun\\,yr^{-1}$, to prevent meridional circulation from canceling the effect of accretion.  The existence of old ($\\approx 1$~Gyr) is not a likely outcome of the present models for accretion/diffusion with or without deep mixing. It is argued that \\lboo\\ stars are potentially very good diagnostics of mixing mechanisms in moderately fast rotators. ",
        "introduction": "The defining characteristic of \\lboo\\ stars is their peculiar surface composition in which only four elements are roughly solar (C, N, O and S) and all others show a definite trend toward depletion by a factor of up to ten typically \\citep{Heiter02,Solanoetal01}. Yet, despite their apparent metal-poor composition these are population~I stars in which the peculiar chemical composition is only superficial. There are concerns that the class, as it is now defined based on their chemical composition, does not represent an homogeneous population resulting from a common physical process.  For example, it is now clear that there is a wide scatter in the abundances of individual elements from star to star, larger than in chemically normal A-type stars \\citep{Heiter02}. Other potentially conflicting observations and the failure of models to reproduce all the characteristics of \\lboo\\ stars, which will be discussed briefly in this paper, caution us to keep that possibility in mind.  In addition to their chemical signature, this class of stars is also characterized by a limited range in spectral types (early F and A type stars) \\citep{Solanoetal01}. For a long time they were also thought to be strictly limited to young stars (ZAMS or pre-main-sequence) but the evidence now points to ages ranging from the ZAMS to the TAMS \\citep{IlievBarzova95,Ilievetal02}. A majority of \\lboo\\ stars are pulsating stars in the general class of \\dscuti-type variable stars \\citep{Bohlenderetal99}. Such pulsations indicate that the superficial helium abundance is roughly solar, or higher. As a much larger fraction of \\lboo\\ stars than other A-type stars are \\dscuti\\ variables \\citep{Paunzenetal98}, it suggests that the \\lboo\\ phenomenon itself has an effect in their pulsational behavior. Finally, many \\lboo\\ stars show signs of circumstellar matter \\citep{Holwegeretal99}.  Several ideas have been put forward to account for the \\lboo\\ phenomenon. The very peculiar abundances of \\lboo\\ stars cannot be reconciled with the diffusion models that have been so successful for other chemically peculiar A stars such as the FmAm stars \\citep{RMT00}. Nevertheless, as diffusion is an important process in A stars, it is an important feature of two leading models proposed for \\lboo\\ stars, the diffusion/mass loss model \\citep{MichaudCharland86,C93} and the accretion/diffusion model \\citep{VL90,C91}. A third hypothesis calls for a binary merging to account for some fraction of \\lboo\\ stars \\citep{Andrievsky97}. It also has been shown that a small  number of \\lboo\\ stars were in fact unidentified binaries for which the combined spectra were mistakenly interpreted as metal-poor \\citep{Faraggianaetal01}. It is doubtful however that a significant fraction of \\lboo\\ stars suffer from this problem. The pro and cons of these hypothesis have been nicely summarized in \\citet{Solanoetal01}. It is clear that at this moment no model can adequately reproduce all the observed properties of \\lboo\\ stars.  The first challenge facing the models is their inability to produce both young and old \\lboo\\ stars. The diffusion/mass loss and binary merging models can only yield old \\lboo\\ stars. On the other hand the accretion/diffusion model as it stands now can only occur in young stars before the circumstellar disk associated with star formation is dissipated.  Currently, the most favored model is the accretion/diffusion model but new observational data has raised questions about it. \\citet{Heiteretal02} have shown that the composition of \\lboo\\ stars might not be as consistent with circumstellar gas has thought previously. They have also raised the question as to why some stars which have circumstellar disks with ongoing accretion and are similar to \\lboo\\ stars do not have peculiar compositions.  This paper concentrates on if and how new insight in the depth of mixing in slowly rotating F and A type stars \\citep{RMT00,RMR01} might affect the accretion model. Those models, which will be discussed further in Section~\\ref{sec:accdiff}, have shown that the mixing in slowly rotating stars extends substantially deeper than the base of the H-He convection zone in A and B~stars. The depth of mixing is of crucial importance in determining the timescales for the formation and persistence of superficial abundance peculiarities in stars such as the \\lboo. Therefore, as a first step in the computation of more complete models for \\lboo\\ stars including as much of the physics of chemical evolution in stars as possible, the possible effect of deep mixing on timescales and accretion rates in models of accreting \\lboo\\ stars will be discussed here. A brief overview of the accretion/diffusion model in the context of deeper mixing will be presented first. Simple models will then be used to estimate timescales for chemical evolution. ",
        "conclusions": "It has been shown that indications of deep mixing in F, A and B stars \\citep{RMT00} can have a significant effect on the accretion/diffusion model for \\lboo\\ stars, leading to larger predicted accretion rates or longer timescales for the formation of the requisite surface composition. This assumes that the mixing in slowly rotating stars is similarly active in faster rotators. Such an extrapolation, as is done here, is still founded on circumstantial evidence and is subject to confirmation.  Nevertheless, deep mixing in \\lboo\\ stars raises intriguing possibilities regarding the few but important points of contention between the standard accretion/diffusion model \\citep{TC93} and the observations \\citep{Solanoetal01,Heiteretal02}. One of the most difficult problem facing the accretion model is the existence of old \\lboo\\ stars. In A type stars, circumstellar disks are not expected to persist more than a couple of hundred Myrs \\citep{MeyerBeckwith00} which is far less than the oldest \\lboo. If one assumes that \\lboo\\ stars are mixed to a depth of $10^{-6}$ in fractional mass, as argued here, and that the accretion rate was high enough early on to ensure the observed abundances reflect those of the accreted matter, then the larger timescale for the evolution of the surface abundance might provide a way to explain older \\lboo\\ stars.  With a standard SMZ only a deep as the SCZ, the time needed to erase the \\lboo\\ signature is of the order of 1~Myr \\citep{TC93}. The timescale in the case of deep mixing will be increased by a factor of some few hundreds because of the increase in the mass of the SMZ, but this is mitigated by the relative increase in flux due to meridional circulation. The net effect for a mixed mass of $10^{-6}$~M$_\\star$ is an increase of the timescale from by a factor of 5 only.  A complicating factor for the accretion scenario is that not only is it necessary to dramatically increase the amount of gas accreted on the star in order to impart the composition of dust-depleted circumstellar gas to the SMZ, which may be accounted by much larger accretion rates on the pre-main-sequence, but much larger ongoing accretion rates are necessary to sustain the abundance peculiarities if the SMZ is as deep as suggested here. Fig.~\\ref{fig:mloss}, shows that an accretion rate of $10^{-12}$ to $10^{-11}$~M$_\\sun\\,yr^{-1}$ is necessary to just balance the flux of particles entering or leaving the SMZ at its base from diffusion and meridional circulation. Such large rates may not problematic as they have been claimed in $\\beta$~Pictoris \\citep{Beustetal96}. It might however raise questions as to whether the amount of circumstellar matter required to provide such large rates could remain unseen as is the case in many \\lboo\\ stars. Still, if one assumes that it is the case and that the necessary accretion is ongoing as long as the circumstellar disk is present, one would still not expect \\lboo\\ stars as old as 1~Gyr.  We have restricted ourselves to rotational velocities of 100~km\\,s$^{-1}$ or lower because of the limitations of the formalism for meridional circulation used here. \\lboo\\ stars can rotate at a much faster rate, as much as 250~km\\,s$^{-1}$ \\citep{Paunzen01}. In such stars, the meridional circulation would dominate a given accretion rate for much shallower SMZs. The accretion rates required to establish the \\lboo\\ signature could then be an order of magnitude larger, or more, than those found for the models discussed here.  Finally, as \\lboo\\ stars often are pulsating stars, it is tantalizing to imagine that there might be a seismic signature of the depth of the mixing considering that the abundance of most metals is completely different in the ``metal opacity bump'' depending on the models discussed here. As the metals play a role in driving pulsations and determining the structure of the envelope in these stars, accretion with deep mixing might yield an observable signature in either which modes become overstable or in shifts in the frequencies of modes of pulsations with respect to standard models for \\dscuti\\ stars. Preliminary models have shown shifts in frequencies by as much as 10 to 30\\% \\citep{T00} but a seismic test for mixing in \\lboo\\ stars, only possible with reliable mode identification, remains out of our reach at this point in time.  \\lboo\\ stars are perhaps the best candidates to provide constraints on mixing mechanisms in moderately rapidly rotating early type stars for which more standard diagnostics such as lithium abundances are not available. The major observational effort spent on these stars in recent years and their inclusion as main targets for planned asteroseismology experiments make a parallel theoretical effort necessary. Only more sophisticated calculations of evolving A stars with accretion and a more precise treatment of rotational mixing and circulation will determine if the accretion/diffusion model can be reconciled with the challenges that the recent observational studies of these stars have uncovered."
    },
    "0206/astro-ph0206434_arXiv.txt": {
        "abstract": "We present the current status of cosmic shear studies and their implications on cosmological models. Theoretical expectations and observational results are discussed in the framework of standard cosmology and CDM scenarios. The potentials of the next generation cosmic shear surveys are discussed. ",
        "introduction": "The gravitational deflection of light beams by large scale structures of the universe (cosmological lensing) amplifies and modifies the shape of distant galaxies and quasars. Magnification produces correlation between the density of foreground lenses and  the apparent luminosity of distant galaxies or quasars (magnification bias), whereas distortion induces a correlation of ellipticity distribution of lensed galaxies (cosmic shear). In both cases, the properties of cosmological lensing signals probe the matter content  and the geometry of universe and how perturbations grew and clustered during the past Gigayears. \\\\ Albeit difficult to detect, the recent cosmic shear detections claimed by several groups demonstrate that it is no longer a technical challenge. It is therefore possible to study the universe through a new window which directly probes dark matter instead of light and allows cosmologists to measure  cosmological parameters and dark matter power spectrum from weak gravitational distortion. ",
        "conclusions": ""
    },
    "0206/astro-ph0206352_arXiv.txt": {
        "abstract": "Isolated low-mass stars are formed in dense cores of molecular clouds. In the standard picture, the cores are envisioned to condense out of strongly magnetized clouds through ambipolar diffusion. Most previous calculations based on this scenario are limited to axisymmetric cloud evolution leading to a single core, which collapses to form an isolated star or stellar system at the center. These calculations are here extended to the nonaxisymmetric case under thin-disk approximation, which allows for a detailed investigation into the process of fragmentation, fundamental to binary, multiple system, and cluster formation. We have shown previously that initially axisymmetric, magnetically subcritical clouds with an $m=2$ density perturbation of modest fractional amplitude ($\\sim 5\\%$) can develop highly elongated bars, which facilitate binary and multiple system formation. In this paper, we show that in the presence of higher order ($m\\ge 3$) perturbations of similar amplitude such clouds are capable of breaking up into a set of discrete dense cores. These multiple cores are magnetically supercritical. They are expected to collapse into single stars or stellar systems individually and, collectively, to form a small stellar group. Our calculations demonstrate that the standard scenario for single star formation involving magnetically subcritical clouds and ambipolar diffusion can readily produce more than one star, provided that the cloud mass is well above the Jeans limit and relatively uniformly distributed. The fragments develop in the central part of the cloud, after the region has become magnetically supercritical but before rapid collapse sets in. It is enhanced by the flattening of mass distribution along the field lines and by the magnetic tension force, which is strong enough during the subcritical-to-supercritical transition to balance out the gravity to a large extent and thus lengthen the time for perturbations to grow and fragments to separate out from the background. ",
        "introduction": "\\label{sec:introduction}  Dense cores of molecular clouds play a pivotal role in star formation (Myers 1999). They provide a crucial link between the molecular clouds and the stars formed in them. In the standard picture for isolated low-mass star formation, the cores are envisioned to gradually condense out of a magnetically subcritical background cloud, through ambipolar diffusion (Shu, Adams \\& Lizano 1987; Mouschovias \\& Ciolek 1999; see Nakano 1998 and Myers 1999 for an alternative view involving turbulence decay). Detailed calculations based on this scenario have been carried out by many authors (e.g., Nakano 1979; Lizano \\& Shu 1989; Ciolek \\& Mouschovias 1993; Basu \\& Mouschovias 1994). It has been established that prior to star formation dense cores have (1) a central region of flat density distribution surrounded by a roughly $r^{-2}$ envelope, (2) an infall speed over an extended region that could be a significant fraction of the isothermal sound speed, and (3) a field strength typically half the critical value. All these features are consistent with the observations of L1544 (Tafalla et al. 1998; Ward-Thompson et al. 1999; Crutcher \\& Troland 2000), arguably the best studied starless core (e.g., Caselli et al. 2002). Ambipolar diffusion-driven cloud evolution models computed specifically for this source (without explicit treatment of turbulence) can match quantitatively its observed mass distribution, velocity field and magnetic field strength (Ciolek \\& Basu 2000), and plausibly the abundances and spatial distributions of various commonly-observed molecular species, such as CO, N$_2$H$^+$, and CCS, as well (Li et al. 2002). For L1544, the standard scenario appears to provide a reasonable description (see Crutcher et al. 1994 for the successful application of a similar model to the dark cloud B1, for which the field strength is also measured).  Most calculations based on the standard scenario are axisymmetric. As such, they can not directly address the fundamental issue of cloud fragmentation, which lies at the heart of binary, multiple stellar system and cluster formation (Bodenheimer et al. 2000). Indeed, there has been some nagging concern whether a magnetic field strong enough to provide most of the cloud support would at the same time prevent fragmentation (Galli et al. 2001). This is certainly the case if the field is completely frozen in the matter (Nakano 1988). Even a somewhat weaker frozen-in field which does not prevent a cloud from collapsing dynamically may stifle fragmentation (Dorfi 1982; Phillips 1986a,b). However, these results have limited applicability to molecular clouds, which are only lightly ionized and partially coupled to the magnetic field. For a lightly ionized medium of uniform density threaded by a uniform magnetic field, Langer (1978) showed through linear analysis that the criterion for gravitational instability is unaffected by the field, although the growth rate can be. For a strongly magnetized dark cloud, the instability grows on an ambipolar diffusion time scale, which is typically an order of magnitude longer than the dynamic time. Nonlinear developments of this magnetically-mediated gravitational instability in molecular clouds have not been explored in any detail (see Indebetouw \\& Zweibel 2000 for simulations of a related instability). How or even whether they can lead to cloud fragmentation into discrete pieces capable of forming more than one star remains uncertain. It is the focus of our investigation.  There is some indication that a strong magnetic field may actually promote molecular cloud fragmentation. Boss (2002) followed the evolution of a set of initially magnetically supported clouds in three dimensions (3D), taking into account several magnetic effects approximately. He concluded that the cloud fragmentation is enhanced by magnetic fields because the magnetic tension helps to prevent a central density singularity from forming and producing a dominant single object. Li (2001) studied the 1D evolution of a set of flattened, magnetically subcritical clouds assuming axisymmetry and found that either a dense supercritical core or off-centered ring forms as a result of ambipolar diffusion. Nakamura \\& Li (2002) showed through 2D nonaxisymmetric calculations that the core-forming clouds are unstable to the $m=2$ nonaxisymmetric mode, with density perturbations of modest fractional amplitude ($\\sim 5\\%$) growing nonlinearly into bars of a typical aspect ratio $\\sim 2$ during the transition period after the core has just become supercritical (which makes the bar growth possible) but before rapid collapse sets in (which leaves little time for further growth). They found that, by the time the isothermal approximation starts to break down, the elongation has been strongly amplified by the Lin-Mestel-Shu (1965) instability, producing highly elongated bars. These bars are expected to break up gravitationally into pieces during the subsequent adiabatic phase of evolution, which could lead to the formation of binary and multiple stellar systems. The possible bar fragmentation into multiple systems will be explored elsewhere. Here, we concentrate on the evolution of nonaxisymmetrically perturbed, ring-forming clouds which, as we show in the paper, are capable of producing a number of discrete, magnetically supercritical cores while still in the isothermal phase of cloud evolution. Even though we cannot follow the evolution of the cores beyond the isothermal phase, we anticipate each of them to collapse individually into a single star or stellar system and, collectively, to form a stellar group or small cluster.  We describe our formulation of the problem of magnetic cloud evolution and fragmentation, including the governing equations, initial conditions, and numerical methods in \\S~2. This is followed by a set of representative models illustrating the main features of nonaxisymmetric cloud evolution leading to fragmentation (\\S~3). In the last section (\\S~4), we discuss the nature of magnetic cloud fragmentation, comment on the implications of our calculations on stellar group formation, and conclude. ",
        "conclusions": "\\subsection{Nature of Magnetic Cloud Fragmentation}  The beneficial roles of strong, ordered magnetic fields in resolving the thorny ``angular momentum problem'' of star formation and in preventing the overall dynamical collapse of molecular clouds are well known (Shu et al. 1987; Mouschovias \\& Ciolek 1999). Their effects on cloud fragmentation are less explored, and appear more subtle. On the one hand, the magnetic pressure can assist the thermal pressure in erasing density inhomogeneities and thus impede fragmentation. On the other hand, the magnetic tension force dilutes gravity, allowing more time for perturbations to grow and over-dense fragments to separate out (see also Boss 2002). The deleterious effects are brought out clearly in the simulations of Phillips (1986a), where the (frozen-in) magnetic fields are taken to be uniform initially, with zero tension force. Phillips (1986a) followed the collapse of a set of such uniformly magnetized clouds with a range in the ratio of thermal to gravitational energies and degree of magnetization. He found that none of the clouds fragmented, even in the presence of a large, $50\\%$ density perturbation. We believe that the lack of fragmentation is mainly due to the weakness of the magnetic tension force relative to gravity in his models, which does not slow down the nearly free-fall collapse appreciably. The tension force is stronger in the models of Phillips (1986b) where the initial fields are non-uniform. However, nonaxisymmetric perturbations are not imposed on these clouds, and it is not clear whether fragmentation is enhanced by the stronger tension force or not. The beneficial effects of magnetic tension force on fragmentation is illustrated most cleanly with the thin-sheet model of Shu \\& Li (1997), which is threaded everywhere by a (frozen-in) magnetic field of critical strength. In such an idealized cloud, the gravity is exactly balanced by the tension force for arbitrary mass distributions, including those with large numbers of condensations (see Allen \\& Shu 2000 for a toy model of the fragmentation of critically magnetized clouds induced by protostellar winds).  The opposing effects of magnetic fields on cloud fragmentation call for a numerical attack, especially in the presence of ambipolar diffusion, which changes the relative importance of the two with time. Our calculations demonstrate that, despite these extra complications, the ambipolar diffusion-driven fragmentation of magnetically subcritical clouds has much in common with that of non-magnetic clouds (e.g., Boss 1996; Klapp \\& Sigalotti 1998): those clouds containing more Jeans masses and having flatter mass distributions are more susceptible to fragmentation. These trends can be understood qualitatively from the classical Jeans analysis (e.g., Larson 1985) and its extension to lightly ionized, strongly magnetized media (Langer 1978). In both cases, cloud fragmentation involves multiple Jeans masses and is driven by the (magnetically diluted) gravity over the resistance of the (magnetically enhanced) pressure gradient. A major advantage of magnetic clouds over their non-magnetic counterparts is their ability to contain multiple Jeans masses without collapsing promptly. Equally important to fragmentation is the flattening of mass distribution along field lines, which enables self-gravitating off-center pockets to preferentially collapse onto themselves rather than falling toward the center (Larson 1985).  Once the central region of a magnetically supported, multi-Jeans mass cloud has become supercritical, nonaxisymmetric perturbations start to grow. The time for growth is lengthened considerably by the magnetic tension force, which remains strong enough during the initial period of the supercritical phase to balance out the gravity to a large extent. Nakamura \\& Li (2002) demonstrated that m=2 density perturbations of modest fractional amplitude of $5\\%$ can grow nonlinearly into bars of aspect ratio $\\sim 2$ during this period (see \\S~\\ref{results} for more examples). It is also during this period that cloud breakup into multiple cores occurs, as speculated by Li (2001) and shown explicitly in the paper. This critical phase of cloud evolution enables the (initially) magnetically subcritical clouds to either break up directly into pieces and/or become significantly elongated, setting the stage for possible (further) fragmentation at later times. It is absent from the evolution of both the non-magnetic and (strongly) magnetically supercritical clouds. The beneficial effects of a strong magnetic field on fragmentation are illustrated vividly in Fig.~1, which shows that increasing the field strength of a subcritical cloud makes it easier for the cloud to break up into multiple cores.  Fragmentation of strongly magnetized cloud cores has been examined by Boss in a series of papers (see Boss 2002 for the latest). Using a 3D code that treats the magnetic forces and ambipolar diffusion approximately, he finds that prolate magnetic cores tend to produce binaries and oblate cores multiple systems. Our calculations are similar to his oblate core calculations but with a different focus: the formation of multiple dense cores out of a flattened, more diffuse background cloud \\footnote{Even though our discussions of the dimensionless solutions of cloud evolution are centered on relatively low initial densities appropriate for pre-core conditions, it may be possible to reinterpret these solutions using (oblate) dense cores as the starting point (if such cores are magnetically subcritical). In this reinterpretation, the discrete overdense blobs developed toward the end of our calculations would have much higher densities and smaller separations, comparable to the multiple protostellar fragments found by Boss (2002). The fact that these two very different sets of calculations (different geometries, numerical methods and treatments of radiative transfer, magnetic field and ambipolar diffusion) yield qualitatively similar results regarding fragmentation is reassuring.}. We concur with Boss's general conclusion that fragmentation can take place {\\it despite} the presence of significant magnetic fields. Indeed, one may go one step further, and speculate that the fragmentation of relatively quiescent clouds into binaries, multiple systems and small groups occurs largely {\\it because of} the strong magnetic field, although more detailed calculations are needed to back up this viewpoint.  \\subsection{Implications on Small Stellar Group Formation}  Magnetically supported clouds can have masses well above the Jeans limit and a naturally-flattened mass distribution, both of which are conducive to the growth of nonaxisymmetric perturbations as the magnetic support weakens with time due to ambipolar diffusion. Nonlinear growth of perturbations produces two basic types of over-dense structure in the supercritical phase of cloud evolution: single, isolated bars for the $m=2$ mode and, for higher order modes, discrete multiple cores, each of which is also significantly elongated in general. Nakamura \\& Li (2002) have speculated that the elongated dense cores (or bars), with their elongation strongly amplified at higher densities by the Lin-Mestel-Shu (1965) instability, may fragment during the subsequent adiabatic phase of cloud evolution, individually producing perhaps a binary or multiple stellar system. More detailed calculations are needed to put the speculation on a firmer ground. Here, we wish to concentrate on the implications of the multiple cores, whose collapse to form collectively a small group of stars appears more certain. This is because by the end of the computations (some of) the cores have cleanly separated out from the background and from one another, and are well on their way to produce stellar density objects.  Some of the salient features of multiple core formation can be seen from Fig.~2, where the evolution and fragmentation of the benchmark cloud with a slow rotation are displayed. Adopting a fiducial temperature of $10$~K and visual extinction of $A_V=1$, we find that the cloud initially contains a central region of flat surface density distribution of radius $\\sim 1$~pc and mass $\\sim 50$~$M_\\odot$. It takes about 5.2 Myrs for the cloud to reach the supercritical state, and another 2.4 Myrs for the m=4 perturbations to grow into well-defined multiple dense cores. The cores are separated by $\\sim 0.05$~pc and, at the end of the calculation, have peak extinction of over 300. They are embedded within a much larger common envelope of lower visual extinction. The cores are well on their way to dynamic collapse from inside-out, each expected to form a protostar or protostellar system which grows by accreting mass competitively from the envelope. In this particular example, we anticipate the formation of a group of stars within a relatively small region of $0.05$~pc or less, although detailed properties of the group, such as the final masses of individual members and their orbital elements, are unknown. Our calculation sets the stage for future higher resolution calculations capable of following the subsequent core evolution to the formation of point mass (protostar) and beyond, which would help elucidate such important topics as competitive mass accretion and orbital evolution of the (still accreting) protostars (Bonnell et al. 1997). Substantial orbital evolution is expected of the group members, not only because the dense cores---the raw materials out of which they are formed---are not rotationally supported (from collapsing into one another), but also because the stars can interact gravitationally with the envelope (by raising tides) and with one another. The former tends to shrink the stellar orbits, and the latter may lead to the ejection of some (lighter) members while leaving others more tightly bound. The decay of small-number groups may be an important channel for producing binaries and multiple systems (e.g., Sterzik \\& Durisen 1999) and, in the scenario of Reipurth \\& Clarke (2001), be responsible for brown dwarf formation.  How are the magnetically subcritical clouds (or ``clumps'' following the notations of Williams et al. 2000) of multiple Jeans masses capable of fragmentation created in the first place? We do not have a definitive answer, but suspect that it probably involves in one way or another turbulence, which is strong in the low density regions out of which the subcritical clouds or clumps are condensed. The condensation could occur for example through compression in the converging regions of a turbulent flow (e.g., Ostriker et al. 1999), and/or as a result of localized turbulence decay (Myers 1999). Since the turbulence is highly supersonic, the thermal pressure (and thus the Jeans limit) should have little relevance to the process. One therefore expects clouds or clumps of more than one Jeans mass to be a common product, at least for those more massive, self-gravitating ones that are the sites of star formation. Whether such clouds or clumps are magnetized strongly enough to be subcritical or not is unclear. As yet, there is not enough direct Zeeman measurements of field strength in relatively low density regions to provide a firm answer to this crucial question one way or the other (Crutcher 1999). In the the standard scenario of isolated low-mass star formation the clumps are envisioned to be magnetically subcritical (Shu et al. 1987; Mouschovias \\& Ciolek 1999). If this turns out to be the case, then their propensity for containing more than one Jeans mass would make them suitable for group formation.  There is some observational evidence that small groups are a common product of low-mass star formation, in both the so-called ``isolated'' and ``clustered'' modes (Shu et al. 1987). The archetype of the former is the Taurus molecular cloud. Molecular line surveys of this nearby dark cloud in $^{13}$CO (Mizuno et al. 1995; at a spatial resolution of $\\sim 0.1$~pc), C$^{18}$O (Onishi et al. 1996; $\\sim 0.1$~pc) and H$^{13}$CO$^+$ (Onishi et al. 1999; $\\sim 0.01$~pc) have found an apparent threshold in the H$_2$ surface density of $\\sim 8\\times 10^{21}$~cm$^{-2}$ for ongoing star formation, as marked by either dense compact (pre-stellar) H$^{13}$CO$^+$ cores or ``cold'' (in IR colors) young stellar objects (YSOs; Onishi et al. 1998). Interestingly, those star-forming C$^{18}$O cores above the threshold typically contain more than one compact core and/or cold YSO, with the more massive ones associated with a larger number of cold objects. Higher resolution continuum maps have also revealed many examples of dust condensations in close proximity, with separations of order $\\sim 0.1$~pc or less (e.g., Looney, Mundy \\& Welch 2000; Shirley et al. 2000; Motte \\& Andre 2001). These groupings of pre-stellar objects and YSOs in relatively quiescent regions could result from the ambipolar diffusion-driven fragmentation of magnetically supported clouds discussed in this paper. Even more striking examples of such groupings can be found in cluster-forming regions, such as the fragmented ring-like structure in the Serpens cloud core (Williams \\& Myers 2000) and the dozen or so starless fragments in the $\\rho$ Oph B2 core (Motte et al. 1998). Whether they could be produced in the same manner is less certain, complicated by the strong turbulence, which is present in these regions but not treated explicitly in our calculations.  One prediction of our scenario is that the group members would have rather small spread in age even though they are formed through ambipolar diffusion, a process that can take up to 10 Myrs. The reason is that the dense cores are produced in our picture together during the supercritical phase, which lasts for a small fraction of the total time of cloud evolution. Even though the denser of the cores would collapse to form stars first, the spread in stellar age should still be relatively small, perhaps comparable to that of stars formed from the fragmentation of non-magnetic clouds---much less than the spread of those formed from separate subcritical clouds in complete isolation from one another.  \\subsection{Conclusion and Future Work}  To summarize, we have followed numerically the ambipolar diffusion-driven evolution of magnetically supported molecular clouds in the presence of various nonaxisymmetric modes of perturbations. We confirmed our previous finding that the $m=2$ perturbations of modest amplitude grow readily into bars, which have implications for binary and multiple system formation. Clouds perturbed with higher order modes can break up into multiple cores, which should collapse to form collectively small groups of stars. These calculations are the first step toward a comprehensive theory of multiple star formation in strongly magnetized clouds.  Our calculations are limited to the isothermal regime. They will be extended into the adiabatic regime at higher densities, to follow the expected breakup of the (long) bars and multiple cores, which are also significantly elongated in general. A new feature at the higher densities is the magnetic decoupling (e.g., Nishi et al. 1991), which will require a more elaborate treatment of the magnetic coupling than the one we adopted. Ultimately, we would like to continue our calculations into the protostellar accretion phase, hoping to determine the masses of individual stars and their orbits, perhaps with an approximate treatment of the magnetic braking which may significantly affect the angular momentum evolution of the system.   Numerical computations in this work were carried out at the Yukawa Institute Computer Facilities, Kyoto University. F.N. gratefully acknowledges the support of the JSPS Postdoctoral Fellowships for Research Abroad. We thank the referee for a prompt and helpful report."
    },
    "0206/astro-ph0206164_arXiv.txt": {
        "abstract": "We calculate the spectral energy distribution (SED) of electromagnetic radiation and the spectrum of high energy neutrinos from BL Lac objects in the context of the Synchrotron Proton Blazar Model.  In this model, the high energy hump of the SED is due to accelerated protons, while most of the low energy hump is due to synchrotron radiation by co-accelerated electrons. To accelerate protons to sufficiently high energies to produce the high energy hump, rather high magnetic fields are required. Assuming reasonable emission region volumes and Doppler factors, we then find that in low-frequency peaked BL~Lacs (LBLs), which have higher luminosities than high-frequency peaked BL~Lacs (HBLs), there is a significant contribution to the high frequency hump of the SED from pion photoproduction and subsequent cascading, including synchrotron radiation by muons.  In contrast, in HBLs we find that the high frequency hump of the SED is dominated by proton synchrotron radiation.  We are able to model the SED of typical LBLs and HBLs, and to model the famous 1997 flare of Markarian 501.  We also calculate the expected neutrino output of typical BL Lac objects, and estimate the diffuse neutrino intensity due to all BL Lacs.  Because pion photoproduction is inefficient in HBLs, as protons lose energy predominantly by synchrotron radiation, the contribution of LBLs dominates the diffuse neutrino intensity.  We suggest that nearby LBLs may well be observable with future high-sensitivity TeV gamma-ray telescopes. ",
        "introduction": "Blazars are identified as Optically Violent Variable (OVV) quasars (a sub-class of Flat Spectrum Radio Quasars, FSRQ) and BL Lacs which may be low-frequency or high-frequency peaked BL~Lac objects.  Their broad-band spectra consist of two spectral components which appear as broad `humps' in the spectral energy distribution, and are due to emission from a jet oriented at small angle with respect to the line-of-sight.  The low-energy component is generally believed to be synchrotron emission from relativistic electrons, and extends from the radio to UV or X-ray frequencies.  The origin of the high-energy component, from X-ray to $\\gamma$-ray energies, in some cases to several TeV, is still under debate.  Various models have been proposed, with the most popular ones being the ``leptonic models'', where a relativistic electron-positron jet plasma up-scatters low energy photons to high energies via the Inverse Compton effect.  The emission region moves with relativistic velocities along the jet, and the resulting radiation is highly beamed into the observer's line-of-sight.  The target photon field could come either from the accretion disk surrounding the putative black hole (e.g.\\cite{DS93,BMS97}), possibly partly reprocessed by broad-line region (BLR) clouds (e.g.\\cite{sikora94}) or a dusty torus (e.g. \\cite{blaze2000}, \\cite{DoneaProtheroe2002}), or it could be produced by the relativistic $e^\\pm$ population itself, in so-called synchrotron - self Compton (SSC) models \\cite{MGC92,BM96}. In such leptonic models, it seems plausible to assume that SSC radiation dominates in objects with relatively weak accretion disk radiation such as BL~Lac objects, while in FSRQs external photons provide the dominant target field for the Inverse Compton process.  As an alternative to leptonic models, ``hadronic models'' were proposed more than 10 years ago to explain the $\\gamma$-ray emission from blazars~\\cite{MB89,M93}. Recently M\\\"ucke \\& Protheroe \\cite{MP2000,MP2001a} have discussed in detail the various contributing emission processes.  In hadronic models the relativistic jet consists of relativistic proton and electron components, which again move relativistically along the jet. High-energy radiation is produced through photomeson production, and through proton and muon synchrotron radiation, and subsequent synchrotron-pair cascading in the highly magnetized environment.  Again either external (i.e. from an accretion disk and/or IR-torus \\cite{P96}) or internal photon fields (i.e. produced by synchrotron radiation from the co-accelerated electrons) can serve as the target for photopion production.  Gamma-ray loud BL Lac objects are most likely explained by the latter possibility. These models can, in principle, be distinguished from the leptonic models by the observation of high energy neutrinos generated in decay chains of mesons created in the photoproduction interactions (for a recent review see \\cite{LM2000}).  In this paper, we study the properties of the Synchrotron Proton Blazar (SPB) model \\cite{MP2001a}, where the dominant target photon field is produced by directly accelerated electrons that manifests itself in the blazar SED as the synchrotron hump.  A detailed description of the model itself and its implementation as a Monte-Carlo/numerical code has already been given in \\cite{MP2001a}.  Since this model is for objects with a negligible external component of the radiation field in the jet, we apply it only to BL~Lac objects.  It has been successfully demonstrated that this model reproduces well the observed double-humped blazar SED.  The goal of this paper is to present a comprehensive study of the model's parameter space, a promising tool for discriminating between leptonic models and the hadronic SPB-model.  We apply our model to both LBLs and HBLs, and discuss our results in the light of the suggested LBL/HBL continuity.  In Section 2 we give a brief description of the SPB model. The model is applied to HBLs and LBLs to calculate the SEDs in Section 3. We vary the magnetic field strength and the target photon density and study their effect on the resulting cascade spectrum. These results are used to identify the parameter sets within the SPB-model which are typical for HBLs and LBLs in Sect. 3.2.  In Sect. 3.3 we compare the model predictions to observed SEDs from HBLs and LBLs.  One of the most dramatic properties of blazars is their variability, and this issue is addressed by modeling the evolution of the SEDs during outburst and quiescent stages in Sect. 3.4.  The predicted neutrino emission from these sources is calculated in Section 4. Finally, we discuss our results in Sect. 5. ",
        "conclusions": "We have presented a parameter study of the SPB model proposed recently to explain the observed spectral energy distribution of $\\gamma$-ray loud BL Lac Objects, i.e.\\ HBLs and LBLs.  This model needs strong magnetic fields together with proton, muon and pion synchrotron radiation in order to produce the double-humped structure observed during active phases of $\\gamma$-ray emission, and this is the main difference to the original hadronic ``proton initiated cascade'' (PIC) model \\cite{M93,MB89} which resulted in a rather featureless $\\pi^{0,\\pm}$-cascade spectrum.  If LBLs possess denser jet frame synchrotron photon fields than HBLs, i.e.\\ denser target photon fields for $p\\gamma$ interactions and cascading in our model, then we have shown that the high-energy emission in these two types of objects is of different origin.  While the MeV-TeV radiation from HBLs is dominated by proton synchrotron radiation, in LBLs there is a significant contribution from muon synchrotron radiation at GeV-TeV energies in addition to the proton synchrotron radiation which dominates at MeV-energies.  This is caused by the significantly higher pion (and muon) production rate. Consequently the injected proton spectrum is cut off due to pion production losses in LBL-like objects, while in HBL-like objects proton synchrotron radiation is responsible for the cut off in the proton spectrum. These cutoffs directly translate into the observed high energy photon cutoff at the source if the Doppler factor is known.  A further consequence of different intrinsic target photon fields in HBLs and LBLs is a difference in the ratio of the luminosity in the high energy hump to low energy hump in their cascade spectra (due to proton acceleration). Because opacity effects decisively influence this ratio, the cascade spectra of LBLs have in general smaller ratios than that of HBLs. Also, the high energy peak to ``dip'' luminosity appears to be smaller in the denser LBL-like environments than in HBLs. This is a consequence of the higher pion production rate in LBLs in comparison to HBLs which causes the featureless $\\pi$-cascade to become important and fill in the gap between the two humps with proton and muon radiation.  To demonstrate the difference between LBLs and HBLs in the SPB-model we have fitted the average observed SED of PKS~0716+714 and Mkn~421. In doing so, we found that HBLs need acceleration efficiencies of order unity to give high energy hump energies in the TeV-range, whereas for LBLs acceleration efficiencies of $\\sim 10^{-2}$ seem to account for the observations. LBLs may also produce multi-TeV photons at a lower level than HBLs despite their higher bolometric luminosity. The production mechanism is through muon synchrotron radiation that peaks at a higher energy than the synchrotron radiation of the primary protons. We therefore suggest that nearby LBLs be included in the observing source lists of future high-sensitivity Cherenkov-telescopes (see, e.g., \\cite{Krennrich2001}).  Gamma-ray loud BL Lacs are well-known for their flaring activity, and we have modeled the nearly simultaneous observations of the intermediate and flaring stage of the famous 1997 giant outburst of Mkn~501. An increase of the Doppler factor and acceleration efficiency, together with rising proton and electron density (leading to a denser intrinsic synchrotron target photon field) can account for the observations satisfactorily.  Because the observations were only simultaneous on a one-day time scale at most we believe that our time-independent code is suitable for this simulation. At this point we stress the need for time-resolved simultaneous observations to provide further constrains for blazar modeling, which then must be carried out using time-dependent simulation codes.   Although proton and muon synchrotron emission, and their reprocessed radiation, produce a double-humped structure in typical blazar jet environments, namely one at X-ray energies, and another at GeV-TeV energies, the low-energy synchrotron target photon field dominates over the X-ray hump in the cascade spectrum in nearly all cases presented here.  This seems to suggest that the SS-PIC model proposed by \\cite{R99,R00}, where the observed X-ray hump is due to reprocessed proton and muon synchrotron radiation, is constrained to a rather narrow parameter range, and this is shown in \\cite{R00}. The relatively small Doppler factors favoured in \\cite{R00} imply thick target photon fields, and consequently significant reprocessing leading to comparable power at X-ray and sub-TeV energies but making it difficult to explain the high-energy bump to be at multi-TeV energies.   The dominance of proton synchrotron radiation in HBLs has recently been used by \\cite{aha2000} to consider a blazar model where all proton synchrotron photons escape the completely optically thin emission region, and appear as the high-energy hump in the blazar SED.  This occurs in intrinsically thin or extremely low energetic ambient photon fields (i.e. very high Doppler factors $D\\approx 10$--$30$ are necessary), where $p\\gamma$-interactions and cascading can be neglected, and extremely large magnetic fields of $\\sim 100$~G are necessary to fit the observed SED in this case. If the X-ray emission is produced by synchrotron emission in the strong magnetic field from secondary electrons which are the result of $\\gamma\\gamma$-pair production of the TeV synchrotron photons of the primary protons with an ambient external infrared photon field, then the X-ray flare is expected to lag the $\\gamma$-ray flare.  Since acceleration of electrons is faster than of protons, $t'_{\\rm{acc,p}}\\gg t'_{\\rm{acc,e}}\\approx t'_{\\rm{syn,e}}$, typically by a few hours for blazars, $\\gamma$-rays should lag X-rays if the X-rays stem from the primary co-accelerated electrons.  If electrons are accelerated in the same process as the protons, $\\gamma$-rays from the PIC-processes are in competition with photons from the leptonic SSC process. SSC-photons contribute significantly to the escaping radiation if $u'_B \\ll u'_{\\rm{phot}}$, i.e.\\ $u'_{\\rm{phot}} > 10^{13}$eV cm$^{-3} (B'/30$G$)^2$. So far, all known $\\gamma$-ray loud BL Lac objects have an energy density low-frequency hump $u'_{\\rm{phot}} < 10^{12}$eV cm$^{-3}$ for reasonable Doppler factors ($D\\approx 10$) and emission volumes, $R'^3\\approx (10^{15-17})^3$ cm$^3$. We conclude that SSC emission is negligible in HBLs and LBLs if relativistic protons are the main carrier of the dissipated energy in a highly magnetized jet.  To calculate the total jet luminosity $L_{\\rm{jet}}$ measured in the rest frame of the galaxy we follow the procedure given in \\cite{ProtheroeDonea2002}.  Under the assumption that all particles (electrons and protons) are relativistic one obtains \\begin{eqnarray} L_{\\rm jet} &=& 4 p'_p \\Gamma^2 \\beta c \\pi R'^2 \\left[\\chi_p {(\\Gamma-1)\\over\\Gamma} + 1 + {p'_B\\over p'_p}+  {p'_e\\over p'_p} \\right]\\\\ &=& {L^{\\rm high}_{\\rm obs} \\over D^4 \\zeta_p} \\left[\\chi_p {(\\Gamma-1)\\over\\Gamma} + 1 + {p'_B\\over p'_p}+ {\\zeta_p L^{\\rm low}_{\\rm obs} \\over \\zeta_e L^{\\rm high}_{\\rm obs}} \\right] \\end{eqnarray} where $\\Gamma = (1 - \\beta^2)^{-1} \\approx D/2$ is a good approximation to the Lorentz factor of jets closely aligned to the line of sight.  $L^{\\rm low}_{\\rm obs}$ and $L^{\\rm high}_{\\rm obs}$ are the observed bolometric luminosities of the low and high energy component, respectively, and $\\zeta_e\\approx 1$, $\\zeta_p$ are the radiative efficiencies for electrons and protons. $u'_B=3p'_B$ is the magnetic energy density of a tangled magnetic field and \\begin{equation} p'_p = {L^{\\rm high}_{\\rm obs} \\over 4 D^4 \\zeta_p \\Gamma^2 \\beta c \\pi R'^2} \\end{equation} \\begin{equation} p'_e = {L^{\\rm low}_{\\rm obs} \\over 4 D^4 \\zeta_e \\Gamma^2 \\beta c \\pi R'^2} \\end{equation} gives the jet-frame pressure of injected relativistic protons and electrons, respectively, that would apply in the absence of energy loss mechanisms, and \\begin{eqnarray} \\chi_p &=& {3\\over4}\\left({p'_e \\over p'_p} {1 \\over{\\gamma'_e}_1 \\ln({\\gamma'_e}_2/{\\gamma'_e}_1)} + {1 \\over {\\gamma'_p}_1 \\ln({\\gamma'_p}_2/{\\gamma'_p}_1)}\\right)\\\\ &=& {3\\over4}\\left({\\zeta_p S^{\\rm low}_{\\rm obs} \\over \\zeta_e S^{\\rm high}_{\\rm obs}} {1 \\over{\\gamma'_e}_1 \\ln({\\gamma'_e}_2/{\\gamma'_e}_1)} + {1 \\over {\\gamma'_p}_1 \\ln({\\gamma'_p}_2/{\\gamma'_p}_1)}\\right). \\end{eqnarray} (Note the erroneous Eq.28+29 in \\cite{MP2001a}.) Here ${\\gamma'_e}_1$,${\\gamma'_e}_2$ are the lower and upper limit of the injected electron spectrum, respectively. The ratio of the number of electrons to protons on injection is then \\begin{equation} {n'_e \\over n'_p} = {p'_e \\over p'_p}{m_p \\over m_e}{{\\gamma'_p}_1\\ln({\\gamma'_p}_2/{\\gamma'_p}_1) \\over{\\gamma'_e}_1 \\ln({\\gamma'_e}_2/{\\gamma'_e}_1)}. \\end{equation} The resulting jet power for HBLs typically lie around $10^{45}$erg/s, for LBLs typically $10^{47}$erg/s, higher than in leptonic SSC models but consistent with estimated upper limits for BL Lacs \\cite{CelottiPadovaniGhisellini97}. Injected $n'_e/n'_p$-ratios are typically $10^{-3}-1$ for HBLs and approximately unity for LBLs assuming ${\\gamma'_e}_1 =2$, i.e. most relativistic electrons responsible for the low-energy hump in the SED would be primaries, co-accelerated with the protons.  The addition of cold electrons possibly needed for charge neutrality in HBLs if $n'_e/n'_p<1$ would add little to the estimated jet power.  A caveat in hadronic models is that most processes are rather slow in comparison to leptonic interactions.  Indeed, if pion production dominates the loss processes, variability time scales below $t_{\\rm{var}} \\approx 10^3 (10/D) (30\\mbox{ G}/B')^2$s at the highest photon energies would not be expected.  This limit is based on $u'_B=u'_{\\rm{phot}}$ and $\\eta=1$, and assuming that the size of the emission region does not constrain the variability time scale.  For HBLs proton synchrotron radiation dominates the loss processes. Hence, the smallest variability time scale (again provided $R'$ does not determine $t_{\\rm{var}}$) depends on the Doppler factor, magnetic field and $\\eta$, which in turn determines essentially the high energy photon turnover $\\epsilon_{\\rm{TeV}}$ (in TeV): $t_{\\rm{var}} \\approx 10^4 [(D/10) \\epsilon_{\\rm{TeV}} (B'/30\\mbox{ G})]^{-0.5}$sec with $\\epsilon_{\\rm{TeV}} \\leq 1.1 (D/10)$ TeV corresponding to $\\eta\\leq 1$.  Thus, for extremely high magnetic fields and/or Doppler factors variability on sub-hour time scales can be reached.  The basic difference between the leptonic SSC model and our presented hadronic model is the content of the jet: while leptonic models work with a relativistic electron/positron plasma, our model considers a relativistic electron/proton jet. For fitting the observed SEDs leptonic models need significantly smaller magnetic field values (e.g. $B'=0.497$~G \\cite{Gh98} or 0.8~G \\cite{pian98} for Mkn~501, $B'=0.46$~G for PKS~0716+714 and $B'=0.093$~G for Mkn~421 \\cite{Gh98}) while the size of the emission region and Doppler factor are comparable to the values used in this model (e.g. \\cite{Gh98} gives $R'=10^{16}$cm for Mkn~501 and Mkn~421, and $R'=5\\times 10^{16}$cm for PKS~0716+714, $D$=10, 12 and 15 for Mkn~501, Mkn~421 and PKS~0716+714, respectively, \\cite{pian98} gives $R'=5\\times 10^{15}$cm and $D=15$ for the flaring state of Mkn~501). As a consequence, in leptonic models the particle energy density is often significantly higher than the magnetic field energy density.  If the magnetic field component along the line of sight is much stronger in hadronic models than in leptonic ones, the rotation measure $RM$ on a length scale of the emission region ($\\sim 10^{15}$cm) may provide a tool to distinguish between the competing models. Observations of spatial and temporal variability of the $RM$ in the central parsecs of AGN suggest that the measured $RM$ is indeed intrinsic to the source and no foreground effect (e.g. \\cite{zavala01}). The observed Faraday rotation may therefore serve as a probe of the magnetic field weighted by the electron density in the so-called Faraday screen along the line of sight on the observed length scale. In radio galaxies and quasars the Faraday screen is often considered to be the narrow (NLR) or broad line region (BLR) (e.g. \\cite{taylor}, \\cite{zavala01}), and electron densities are derived from the NLR/BLR optical line strengths. According to the unified scheme the proposed picture for BL Lac Objects consists of a relativistic jet that evacuates a cone through the ionized gas in the nuclear region such that cores of BL Lacs are not viewed through a dense Faraday screen, and lower $RM$-values are therefore expected from BL Lac Objects (e.g. \\cite{taylor}). To date, $RM$ measurements from AGN exist only on kpc-pc scales. E.g., for BL Lacs \\cite{zavala02} recently found $RM\\sim$ several 100 rad~m$^{-2}$ on 1-50pc scales. Assuming $N_e\\approx 10-100$cm$^{-3}$ these values fit to the strong magnetic fields in hadronic models if the field decays along the jet as $R^{-1.5\\ldots -2}$. High-resolution $RM$ observations on the central $10^{-3}$~pc scales and a definite identification of the Faraday screen in BL Lacs are needed to clearly constrain the magnetic field in the gamma ray emission region.   In contrast to leptonic models, models involving pion production inevitably predict neutrino emission due to the decay of charged mesons.  In the present work, we predict the neutrino output of a typical LBL, PKS~0716+714, and a typical HBL, Mkn~421. If LBLs possess intrinsically denser target photon fields than HBLs, then within the SPB model, higher meson production rates are expected in LBLs, leading to a higher neutrino production rate. The diffuse neutrino background will therefore be dominated by LBLs, unless HBLs turn out to be significantly more numerous than LBLs.  To estimate the diffuse neutrino flux we must know the luminosity functions of HBLs and LBLs, and the neutrino energy spectrum of HBLs and LBLs as a function of their luminosity.  Because of the uncertainties in the BL Lac luminosity function and the conversion from low energy peak to 5 GHz-luminosity, the diffuse neutrino flux can only be predicted within a large uncertainty of more than three orders of magnitude. To reduce this uncertainty, it would be helpful to have a luminosity function for $\\nu L_{\\nu}^{\\rm{peak}}$ or the bolometric luminosity of the low-energy hump."
    },
    "0206/astro-ph0206487_arXiv.txt": {
        "abstract": "\\noindent We report on a search for electromagnetic and/or hadronic showers (\\textit{cascades}) induced by a diffuse flux of neutrinos with energies between 5~TeV and 300~TeV from extraterrestrial sources.  Cascades may be produced by matter interactions of all flavors of neutrinos, and contained cascades have better energy resolution and afford better background rejection than through-going $\\nu_\\mu$-induced muons.  Data taken in 1997 with the AMANDA detector were searched for events with a high-energy cascade-like signature. The observed events are consistent with expected backgrounds from atmospheric neutrinos and catastrophic energy losses from atmospheric muons. Effective volumes for all flavors of neutrinos, which allow the calculation of limits for any neutrino flux model, are presented. The limit on cascades from a diffuse flux of $\\nu_e+\\nu_\\mu+\\nu_\\tau+\\overline{\\nu}_e+\\overline{\\nu}_\\mu+\\overline{\\nu}_\\tau$ is $E^2 \\frac{d\\Phi}{dE} < 9.8 \\times 10^{-6}\\;\\mathrm{GeV\\,cm^{-2}\\,s^{-1}\\,sr^{-1}}$, assuming a neutrino flavor flux ratio of 1:1:1 at the detector.  The limit on cascades from a diffuse flux of $\\nu_e+\\overline{\\nu}_e$ is $E^2 \\frac{d\\Phi}{dE} < 6.5 \\times 10^{-6}\\;\\mathrm{ GeV\\,cm^{-2}\\,s^{-1}\\,sr^{-1}}$, independent of the assumed neutrino flavor flux ratio. ",
        "introduction": "Neutrinos interact principally via the weak force, posing a detection challenge for neutrino telescopes but bestowing a valuable advantage on the field of neutrino astronomy: neutrino fluxes from astronomical sources are essentially unattenuated even over cosmological distances. In contrast, high-energy gamma rays are absorbed and/or scattered by intervening matter and photons, and high-energy cosmic-rays are deflected by galactic and intergalactic magnetic fields except at the highest energies ($>10^{19}$~eV).  We present a search for the fully-reconstructed light patterns created by electromagnetic or hadronic showers (\\textit{cascades}) resulting from a diffuse flux of high-energy extraterrestrial neutrinos.  We use data collected in 1997 from the Antarctic Muon and Neutrino Detector Array (AMANDA) for this purpose. Demonstrating $\\nu$-induced cascade sensitivity is an important step for neutrino astronomy because the cascade channel probes all neutrino flavors, whereas the muon channel is primarily sensitive to charged current $\\nu_\\mu$ and $\\overline{\\nu}_\\mu$ interactions.  This is particularly relevant in view of the emerging understanding of neutrino oscillations~\\cite{sno-cc,sno-nc,sno-dn,superk-atm}, in which the flux of $\\nu_\\mu$ would be reduced by oscillations.  (The detection of high-energy atmospheric muon neutrinos by AMANDA has been demonstrated by the full reconstruction of Cherenkov light patterns produced by up-going muons~\\cite{ama:nature-pub,amanda:atm-b10,amanda:wimps-b10}.) Cascades also boast more accurate energy measurement and better separation from background, although they suffer from worse angular resolution and reduced effective volume relative to muons. Importantly, it is straightforward to calibrate the cascade response of neutrino telescopes such as AMANDA at lower energies through use of, e.g., \\textit{in-situ} light sources.  Furthermore, cascades become increasingly easier to identify and reconstruct as detector volumes get larger, so the techniques presented here have relevance for future analyses performed at larger detectors.  Electron neutrinos can produce cascades with no detectable track via the charged current (CC) interaction and all neutrino flavors can produce cascades via the neutral current (NC) interaction.  Cascade-like events are also produced in $\\nu_\\tau$ CC interactions when the resulting $\\tau$ decays into an electron (roughly 18\\% branching ratio) or into mesons (roughly 64\\% branching ratio) and the $\\tau$ energy is below about 100~TeV, at which energy the $\\tau$ decay length is less than 5~m, so that the shower produced by the neutrino interaction and the shower produced by the $\\tau$ decay cannot be spatially resolved by AMANDA. The contribution of $\\nu_\\tau$ to the cascade channel becomes important when flavor oscillations are taken into account for extraterrestrial~\\cite{halzen-saltzberg,reno,beacom} and for atmospheric~\\cite{stanev:prl} $\\nu$-induced cascades. For extraterrestrial sources, current knowledge of neutrino oscillations suggests a detected neutrino flavor flux ratio of $\\nu_e$:$\\nu_\\mu$:$\\nu_\\tau$::1:1:1 following an expected flux ratio of 1:2:0 at the source.  The total light output of an electromagnetic cascade is approximately $10^{8}$ photons/TeV in ice. Hadronic cascades have a light yield about 20\\% lower~\\cite{wiebusch:phd}. An electromagnetic cascade develops in a cylinder of about 10-15~cm in radius (\\textit{Moli\\`ere radius}) and several meters in length (about 8.5~m from the vertex of a 100~TeV cascade, essentially all charged particles are below the critical energy).  Hadronic cascades have longer longitudinal developments and larger Moli\\`ere radii. As a sparsely instrumented detector, AMANDA is insensitive to the topological differences between electromagnetic and hadronic cascades.  Since the NC interaction has a lower cross section and results in a deposition of less energy than the CC interaction, and since we assume a steeply falling neutrino energy spectrum, at any given energy a very small fraction of the $\\nu_e$ events are due to NC interactions.  Hence their impact on the cascade energy resolution is small, and the energy spectrum of reconstructed cascades closely follows that of the CC $\\nu_e$ energy spectrum.  In this paper, we present limits on the diffuse fluxes of ($\\nu_e+\\nu_\\mu+\\nu_\\tau+\\overline{\\nu}_e+\\overline{\\nu}_\\mu+\\overline{\\nu}_\\tau$) and ($\\nu_e+\\overline{\\nu}_e$), assuming a customary $E^{-2}$ power law spectrum at the source.  These limits are based on the observation of no events consistent with a diffuse flux of high-energy extraterrestrial neutrinos. We also present effective volumes for all neutrino flavors to facilitate the calculation of a limit for any flux model.  (A search for up-going muons produced by a $\\nu_\\mu$ extraterrestrial diffuse flux is presently being conducted and preliminary results have been reported in~\\cite{icrc01:dif}.) ",
        "conclusions": "High-energy neutrino-induced cascades have been searched for in the data collected by AMANDA-B10 in 1997. Detailed event reconstruction was performed.  Using \\textit{in-situ} light sources and atmospheric muon catastrophic energy losses, the sensitivity of the detector to high-energy cascades has been demonstrated.  No evidence for the existence of a diffuse flux of neutrinos producing cascade signatures has been found. Effective volumes as a function of energy and zenith angle for all neutrino flavors have been presented. The effective volumes allow the calculation of limits for any predicted neutrino flux model. The limit on cascades from a diffuse flux of $\\nu_e+\\nu_\\mu+\\nu_\\tau+\\overline{\\nu}_e+\\overline{\\nu}_\\mu+\\overline{\\nu}_\\tau$ is $E^2 \\frac{d\\Phi}{dE} < 9.8 \\times 10^{-6}\\;\\mathrm{GeV\\,cm^{-2}\\,s^{-1}\\,sr^{-1}}$ assuming a neutrino flavor flux ratio of 1:1:1 at the detector.  The limit on cascades from a diffuse flux of $\\nu_e+\\overline{\\nu}_e$ is $E^2 \\frac{d\\Phi}{dE} < 6.5 \\times 10^{-6}\\;\\mathrm{ GeV\\,cm^{-2}\\,s^{-1}\\,sr^{-1}}$, independent of the assumed neutrino flux ratio.  The limits are valid for neutrino fluxes in the energy range of 5~TeV to 300~TeV.   \\begin{table*} \\footnotesize{ \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|} \\hline &                        & 3.0-10.0~TeV  & 10.0-30~TeV   & 30-100~TeV    & 100-300~TeV\\\\ \\hline \\hline & $-1<\\cos\\theta<-0.6$   & $0.80\\pm0.05$ & $1.85\\pm0.10$ & $1.87\\pm0.15$ & $1.37\\pm0.20$ \\\\ $\\nu_e$ & $-0.6<\\cos\\theta<-0.2$ & $0.40\\pm0.03$ & $0.85\\pm0.07$ & $1.10\\pm0.10$ & $0.72\\pm0.10$ \\\\ & $-0.2<\\cos\\theta<0.2$  & $0.08\\pm0.01$ & $0.22\\pm0.02$ & $0.36\\pm0.05$ & $0.31\\pm0.07$ \\\\ \\hline & $-1<\\cos\\theta<-0.6$   & $0.82\\pm0.05$ & $1.67\\pm0.12$  & $1.85\\pm0.10$ & $1.60\\pm0.15$ \\\\ $\\overline{\\nu}_e$ & $-0.6<\\cos\\theta<-0.2$ & $0.42\\pm0.03$ & $0.77\\pm 0.07$ & $0.92\\pm0.07$ & $0.74\\pm0.10$   \\\\ & $-0.2<\\cos\\theta<0.2$  & $0.09\\pm0.01$ & $0.20\\pm0.02$  & $0.35\\pm0.05$ & $0.30\\pm0.07$\\\\ \\hline & $-1<\\cos\\theta<-0.6$   & $0.08\\pm0.02$ & $0.35\\pm0.05$ & $0.87\\pm0.1$  & $1.27\\pm0.15$ \\\\ $\\nu_\\mu$ & $-0.6<\\cos\\theta<-0.2$ & $0.05\\pm0.01$ & $0.25\\pm0.03$ & $0.70\\pm0.10$ & $1.60\\pm0.10$ \\\\ & $-0.2<\\cos\\theta<0.2$  & $-$           & $-$           & $-$           & $0.05\\pm0.01$\\\\ \\hline & $-1<\\cos\\theta<-0.6$   & $0.12\\pm0.02$ & $0.34\\pm0.05$ & $0.70\\pm0.05$ & $1.17\\pm0.15$\\\\ $\\overline{\\nu}_\\mu$ & $-0.6<\\cos\\theta<-0.2$ & $0.05\\pm0.01$ & $0.25\\pm0.03$ & $0.70\\pm0.1$  & $0.14\\pm0.01$ \\\\ & $-0.2<\\cos\\theta<0.2$  & $-$           & $-$           & $-$           & $0.03\\pm0.01$\\\\ \\hline & $-1<\\cos\\theta<-0.6$   & $0.35\\pm0.05$ & $1.10\\pm0.10$ & $1.85\\pm0.15$ & $1.35\\pm0.20$ \\\\ $\\nu_\\tau$ & $-0.6<\\cos\\theta<-0.2$ & $0.15\\pm0.03$ & $0.50\\pm0.05$ & $0.85\\pm0.10$ & $1.05\\pm0.10$\\\\ & $-0.2<\\cos\\theta<0.2$  & $0.04\\pm0.01$ & $0.10\\pm0.02$ & $0.23\\pm0.05$ & $0.32\\pm0.07$\\\\ \\hline & $-1<\\cos\\theta<-0.6$   & $0.35\\pm0.05$ & $1.15\\pm0.10$ & $1.65\\pm0.10$ & $1.50\\pm0.15$\\\\ $\\overline{\\nu}_\\tau$ & $-0.6<\\cos\\theta<-0.2$ & $0.15\\pm0.03$ & $0.45\\pm0.05$ & $0.80\\pm0.10$ & $1.20\\pm0.10$ \\\\ & $-0.2<\\cos\\theta<0.2$  & $0.06\\pm0.01$ & $0.12\\pm0.02$ & $0.22\\pm0.04$ & $0.31\\pm0.06$\\\\ \\hline \\end{tabular} \\end{center} \\caption{\\label{table:effec_vol} Effective volume, in units of $10^{-3}$km$^3$, for all neutrino flavors as a function of energy and zenith angle after all the selection criteria have been applied. Uncertainties are statistical only.}} \\end{table*}"
    },
    "0206/astro-ph0206214_arXiv.txt": {
        "abstract": "We present an absolute parallax and relative proper motion for the fundamental distance scale calibrator, $\\delta$ Cep. We obtain these with astrometric data from FGS 3, a white-light interferometer on {\\it HST}. Utilizing spectrophotometric estimates of the absolute parallaxes of our astrometric reference stars and constraining $\\delta$ Cep and reference star HD 213307 to belong to the same association (Cep OB6, de Zeeuw et al. 1999), we find $\\pi_{abs}  =  3.66 \\pm 0.15$ mas. The larger than typical astrometric residuals for the nearby astrometric reference star HD 213307 are found to satisfy Keplerian motion with P = 1.07 $ \\pm $ 0.02 years, a perturbation and period that could be due to a F0V companion $\\sim7$ mas distant from and $\\sim$4 magnitudes fainter than the primary. Spectral classifications and VRIJHKT$_2$M and DDO51 photometry of the astrometric reference frame surrounding $\\delta$ Cep indicate that field extinction is high and variable along this line of sight. However the extinction suffered by the reference star nearest (in angular separation and distance) to $\\delta$ Cep, HD 213307, is lower and nearly the same as for $\\delta$ Cep. Correcting for color differences, we find $<$A$_V>$ = 0.23 $ \\pm $ 0.03 for $\\delta$ Cep, hence, an absolute magnitude M$_V  =  -3.47  \\pm 0.10$. Adopting an average V magnitude, $<$V$>$ =  15.03 $\\pm$ 0.03, for Cepheids with log P  =  0.73 in the LMC from Udalski et al. (1999), we find a V-band distance modulus for the LMC, m-M $ =  18.50 \\pm 0.13$ or, $18.58 \\pm 0.15$, where the latter value results from a highly uncertain metallicity correction (Freedman et al. 2001). These agree with our previous RR Lyr {\\it HST} parallax-based determination of the distance modulus of the LMC. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206022_arXiv.txt": {
        "abstract": "{We obtained deep H- and K-band images of DENIS-P~J104814$-$395606 using SofI and the speckle camera SHARP-I at the ESO-3.5m-NTT as well as QUIRC at the Mauna Kea 2.2m telescope between December 2000 and June 2001. The target was recently discovered as nearby M9-dwarf among DENIS sources (Delfosse et al. 2001). We detect parallactic motion on our images and determine the distance to be $4.6 \\pm 0.3$ pc, more precise than previously known. From the available colors, the distance, and the spectral type, we conclude from theoretical models that the star has a mass of $\\sim 0.075$ to $0.09~$M$_{\\odot}$ and an age of $\\sim 1$ to 2 Gyrs. We also obtained H- and K-band spectra of this star with ISAAC at the VLT. A faint companion candidate is detected 6$^{\\prime \\prime}$ NNW of the star, which is $6.4 \\pm 0.5$ mag fainter in H. However, according to another image taken several month later, the companion candidate is not co-moving with the M9-dwarf. Instead, it is a non-moving background object. Limits for undetected companion candidates are such that we can exclude any stellar companions outside of $\\sim 0.25^{\\prime \\prime}$ (1 AU), any brown dwarf companions (above the deuterium burning mass limit) outside of $\\sim 2^{\\prime \\prime}$ (9 AU), and also any companion down to $\\sim 40$~M$_{\\rm jup}$ with $\\ge 0.15^{\\prime \\prime}$ (0.7 AU) separation, all calculated for an age of 2 Gyrs. Our observations show that direct imaging of sub-stellar companions near the deuterium burning mass limit in orbit around nearby ultra-cool dwarfs is possible, even with separations that are smaller than the semi-major axis of the outermost planet in our solar system, namely a few tens of AU. ",
        "introduction": "The object DENIS-P~J104814$-$395606 (hereafter Denis~1048$-$39) was discovered as very red source in the DEep Near-Infrared Survey (DENIS, Epchtein 1997) by Delfosse et al. (2001): I=$12.67 \\pm 0.05$, J=$9.59 \\pm 0.05$, and K$_{\\rm s}=8.58 \\pm 0.05$ mag. Together with its USNO B- and R-band colors (B=$19.0 \\pm 0.3$ and R=$15.7 \\pm 0.2$ mag), it was classified as late M-type star; a follow-up spectrum taken with Keck showed the spectral type to be M9; from the previous USNO and DENIS imaging, a high proper motion of $1.516 \\pm 0.012^{\\prime \\prime}$ per year to the SE was determined; from its magnitudes and spectral type, a distance of $4.1 \\pm 0.6$ pc was determined; its galactic UVW velocities are typical for the thin disk (Delfosse et al. 2001). Deacon \\& Hambly (2001) obtained its trigonometric parallax from archived plates to be $192 \\pm 37$ mas, i.e. a distance of $5.2 ^{+ 1.2} _{-0.8}$ pc, still consistent with the photometric distance estimate from Delfosse et al. (2001) for a single or binary star.  Because of its small distance and intrinsically faint magnitude, such an object is a promising target for direct imaging detection of sub-stellar companions, both brown dwarfs and giant planets. Due to the problem of dynamic range, namely that sub-stellar objects are too faint and too close to their primary stars, it is difficult to directly detect such objects in orbit around a normal star. One way around this problem is to observe intrinsically faint objects like so-called ultra-cool dwarfs (i.e. dwarfs with spectral type M6 or later, including brown dwarfs). If they are also very nearby, one can probe separations which are of the order of the semi-major axes of the giant planets of our own solar system.  So far, only a few brown dwarf companion candidates detected directly in orbit around normal stars are confirmed by both spectroscopy and proper motion, the first of which was Gl~229~B (Nakajima et al. 1995, Oppenheimer et al. 1995), and the youngest of which is TWA-5~B (Lowrance et al. 1999, Neuh\\\"auser et al. 2000). The first four confirmed brown dwarf companions\\footnote{in addition to Gl~229~B and TWA-5~B mentioned above, there are G~196-3~B (Rebolo et al. 1998) and GD~570~D (Burgasser et al. 2000)} all orbit M-type dwarf stars. Only the fifth such companion, HR~7329~B, orbits an A0-type star (Lowrance et al. 2000, Guenther et al. 2001). This may indicate that nature prefers binaries with not too different masses, another motivation to search for sub-stellar companions around M-type stars; however, it could also be a selection effect, because companions with very different masses and magnitudes are much more difficult to detect.  In Sect. 2, we show our first two deep infrared (IR) imaging observation of Denis~1048$-$39. Follow-up H- and K-band spectroscopy is presented in Sect. 3. Two more follow-up images of Denis~1048$-$39 are used to derive the proper motion of the companion candidate in Sect. 4 and the parallax of Denis~1048$-$39 itself in Sect. 5. Then, in Sect. 6, we search for closer companions with the speckle camera SHARP-I. We discuss and summarize our results in Sect. 7. ",
        "conclusions": "A faint, possibly sub-stellar companion candidate 6$^{\\prime \\prime}$ NNW of Denis~1048$-$39 turned out to be a background object. Because several other groups are actively searching for such sub-stellar companions around nearby stars, we think that it is important also to report such a negative result. Because the primary object itself is almost sub-stellar, any real companion which is just a few mag fainter, would be sub-stellar. Because of the high dynamic range achieved with the speckle camera SHARP-I, we can exclude stellar companions at separations above $0.25^{\\prime \\prime}$, which is only 1 AU at 4.6 pc distance, and we can also exclude brown dwarf companions at $\\ge 2^{\\prime \\prime}$ (9 AU), all for an age of 2 Gyrs.  Three epoch imaging led to the determination of the parallax of Denis~1048$-$39 corresponding to a distance of $4.6 \\pm 0.3$ pc. This distance estimate is consistent with both the photometric distance of $4.1 \\pm 0.6$ pc given in Delfosse et al. (2001) in case that the object is a single main sequence M9-dwarf as well as with the trigonometric parallax measurement (from archived plates) by Deacon \\& Hambly (2001). At this distance, its apparent magnitudes (assuming no interstellar reddening) the following absolute magnitudes, rarely available for an M9-dwarf: M$_{\\rm B}=20.69 \\pm 0.33$, M$_{\\rm R}=17.39 \\pm 0.24$, M$_{\\rm I}=14.36 \\pm 0.14$, M$_{\\rm J}=11.28 \\pm 0.14$, M$_{\\rm H}=11.24 \\pm 0.34$, and M$_{\\rm K}=10.27 \\pm 0.14$ mag.  At the given distance, its colors and spectral type place the object into the H-R diagram, so that we can compare its position with theoretical models (Burrows et al. 1997, Chabrier et al. 2000). The object is a very low-mass star with a mass of $\\sim 0.075$ to $\\sim 0.09$~M$_{\\odot}$ and an age of $\\sim 1$ to 2 Gyrs. There is no star known that is both more nearby and cooler than Denis~1048$-$39. At a distance of $\\sim 4.6$~pc, it is the 40th nearest star.\\footnote{see http://www.chara.gsu.edu/RECONS/TOP100.htm}  Speckle imaging with SHARP has revealed no additional closer companion candidates, so that it seems likely that Denis~1048$-$39 is a single star. However, an even lower mass or even closer companion cannot be excluded.  Our direct imaging observations of the companion candidate at $\\sim 6^{\\prime \\prime}$ separation with an H-band magnitude difference of $\\sim 6.4$ mag shows that such faint possibly sub-stellar companions can be detected directly around nearby ultra-cool dwarfs, in this case at a projected physical separation of only $\\sim 27$~AU, if it were bound. In the near future, this target and other similar nearby ultra-cool dwarfs can be observed with the repaired NICMOS at the HST and CONICA-NAOS at the VLT, which provides an IR wavefront sensor. Then, we can investigate the multiplicity of such objects with even higher dynamic range and higher sensitivity."
    },
    "0206/astro-ph0206304_arXiv.txt": {
        "abstract": "{  In this paper we present a test case for the existence of a core in the density distribution of dark halos around galaxies. DDO 47 has a rotation curve that increases linearly from the first data point, at $300 \\ pc$, up to the last one, at $ 5 \\ kpc$. This profile implies the presence of a (dark) halo with an (approximately) constant density over the region mapped by data. This evidences the inability of standard $\\Lambda$ Cold Dark Matter scenario to account for the dark matter distribution around galaxies, and points toward the existence of an intriguing halo scale--length of homogeneity. This work adds up to the results of Blais-Ouellette et al (2002), Trott \\& Webster (2002), Binney \\& Evans (2002), de Blok \\& Bosma (2002) and  Bottema (2002) in suggesting that CDM theory should incorporate, as an intrinsic property at galactic scales, a ``density core\" feature.} ",
        "introduction": "Rotation curves (RC's) of disk galaxies are the best probes for dark matter (DM) on galactic scale. Although much progress has been made over the past 20 years, it is only very recently that we start to shed light on crucial aspects of the DM {\\it distribution}. Initially, the main focus was on the presence of a dark component; this later shifted to investigating the ratio of dark to visible matter (Salucci and Persic, 1997). Today, the focus is mainly on the actual density profile of dark halos (e.g. Salucci, 2001) A cored distribution, i.e. a  density profile flat out to a radius that is a significant part of the disk size,  has been often adopted (e.g. Carignan \\& Freeman, 1985), although the implications of this distribution appeared only after that  cosmological $N$--body simulations found that Cold Dark Matter (CDM) virialized halos achieve a cuspy density profile (Navarro, Frenk \\& White, 1995, hereafter NFW): \\begin{equation} \\rho_{CDM}(r) =\\frac{\\rho_s}{x(1+x)^2} \\label{rhoCDM} \\end{equation} where $x= r/r_s$, $r_s$ and $\\rho_s$ are the characteristic inner radius and density and the simulations' spatial resolution has  recently  reached  $1/10 r_s$. The halo virial radius $R_{vir }$ is the radius within which the mean halo density is $\\Delta_{vir}(z) $ times the mean cosmic density at that red shift (see Bullock et al., 2001). The virial mass $M_{vir}$ and the corresponding virial velocity are related by: $V_{ vir}^2 \\equiv G M_{vir} / R_{vir}$. The concentration parameter is defined by $c \\equiv R_{vir}/r_s$. With the above definitions, the NFW circular velocity can be written as:  \\begin{equation} V_{CDM}^2(r)= V_{vir}^2 \\frac{c}{A(c)} \\frac {A(x)}{x} \\label{Vh} \\end{equation} where $A(x)= \\ln (1+x) - x/(1+x)$.  Numerical simulations show that in  CDM halos  the virial mass, the virial radius and the concentration are  mutually related: objects considered at $z=0$ within the cosmological scenario with $\\Lambda=0.7$, $\\Omega_0=0.3$ and $h=0.7$ have:  \\begin{equation} c \\simeq 21 \\ \\left ( \\frac{M_{vir}}{10^{11}\\ M_{\\odot}}\\right )^{-0.13} \\\\ R_{vir} \\simeq 120 \\left ( \\frac{M_{vir}}{10^{11}\\ M_{\\odot}}\\right )^{1/3} \\ kpc \\end{equation}  (Wechsler et al., 2002). By applying the above to a reference mass of $5\\times 10^{10} M_\\odot$, reasonable for DDO 47, we get: $c \\simeq 22$ and $r_s \\simeq 4 \\ kpc$, that are very solid guesses in view of the weak mass dependence of these quantities. The available HI data, then, map, exactly the region in which a NFW halo changes its velocity slope from 0.5 to 0. The relationship (3) frames the $\\Lambda$CDM  halo around DDO 47,  but will not enter in the crucial evidence we bring for $\\Lambda$CDM;    the evidence, in fact,   will concern the theory at the most  fundamental level of eq (2).  Discrepancies between the universal profile of CDM and the mass distribution of the dark halo as inferred from the RC has emerged a few years ago (Moore, 1994) At the present, the existence of a crisis for CDM is seriously considered  (Blais-Ouellette et al (2002), Trott \\& Webster (2002), Binney \\& Evans (2002), de Blok \\& Bosma (2002), Bottema (2002)), but the  opposite view is also claimed (van den Bosch et al., 2000, Primack, 2002) and   so as  the view for which  NFW halos fare badly but not worse than other profiles (Jimenez et al., 2002).   Our  previous works approached the issue in two complementary ways: {\\it i)} we derived the accurate mass structure of halos around galaxies and then tried to fit them inside the CDM scenario (Borriello \\& Salucci, 2001)  and {\\it ii)} we tested strategic CDM features by means of appropriate available kinematical data (Salucci, 2001).     The recent HI data (Walter \\& Brinks, 2001, hereafter WB01) and $I$--band surface brightness photometry (Makarova et al., 2002) relative to  the dwarf galaxy DDO 47 give the opportunity to combine these two  approaches,  provide also an exemplar test case for CDM.  In fact: {\\it a)} the RC of DDO 47 extends out to $\\sim 9$ disk scale-lengths, which correspond to 1.3 NFW halo scale--length $ r_s$, at a spatial resolution of $1/7 r_s$ (i.e.  one disk scale-length) {\\it b)} the HI disk surface density  decreases sharply with radius: its flat contribution to the circular velocity does not mimic the solid-body profile of a constant density halo, and therefore does not complicate the mass modeling. {\\it c)} the galaxy is of low luminosity: the content of luminous mass is small with respect to the dark one (e.g. Persic and Salucci,  1988) and consequently easy to take into account.  With these favorable circumstances, we are able to correctly investigate the mass structure of the dark halo around DDO 47, and eventually to discover the inner $r^{-1}$ signature of the NFW universal profile.  Finally,  we do not consider the compression exerted by  baryons when they infall on dark halos, in that this process makes CDM halos density profiles even steeper than the original NFW ones (Blumenthal, 1986).  DDO 47 is assumed at a distance of 4 Mpc, so that $1\"=19.38 pc$ . The crucial results of this paper do not depend on the actual value for the galaxy distance, however, we will discuss the marginal role it plays.  \\begin{figure*} \\center{\\psfig{file=f1.ps,width=140mm,height=140mm}} \\vspace{-9truecm} \\caption{The HI surface density of DDO 47 (points) with the fit used in eq(6) (solid line). Units are $M_{\\odot}/{\\rm pc}^2$ {\\it vs} arcsec } \\label{HI} \\end{figure*} ",
        "conclusions": "We think that $\\Lambda$CDM NFW halos  have  lost, in DDO 47, the last call to represent the dark matter   around galaxies. This galaxy, in fact, in terms of extension, spatial resolution, regularity and smallness of the observational errors of the rotation curve and  in terms of the large dark-to-luminous mass and $r_s$/(spatial resolution of rotation curve)  ratios,   is a perfect laboratory  to detect a NFW halo  by pin-pointing the density slope change, from 0.5 to about  0,  that should occur exactly in the region mapped by data. Saying it plainly,  DM halo density of DDO 47, out to $9\\ r_d$, is instead  fully and uniquely determined by two parameters, a  core density and a core radius,   that are {\\it not even existing} in the gravitational instability/hierarchical clustering Cold Dark Matter scenario.      To reconcile the  $\\Lambda$CDM theory with this evidence is clearly beyond the scope of this paper, let us however  just indicate two possible routes: 1- the (gravitational) physics of  the collapse of  the innermost 10\\%  of the halo mass distribution could be  more complex than that modeled by  current CDM simulations. 2- a (yet) unknown physical process could occur in the innermost $10^{-3}$\\% of the dark halo volume, cutting down the post--collapse DM density by $1-2$ orders of magnitudes."
    },
    "0206/astro-ph0206418_arXiv.txt": {
        "abstract": "{An \\emph{XMM--Newton} study of ultraluminous X--ray sources (ULX) has been performed in a sample of 10 nearby Seyfert galaxies. Eighteen ULX have been found with positional uncertainty of about $4''$. The large collecting area of \\emph{XMM--Newton} makes the statistics sufficient to perform spectral fitting with simple models in 8 cases. The main results of the present minisurvey strengthen the theory that the ULX could be accreting black holes in hard or soft state. In some cases, the contribution of the ULX to the overall X--ray flux appears to be dominant with respect to that of the active nucleus. In addition, 6 ULX present probable counterparts at other wavelengths (optical/infrared, radio). A multiwavelength observing strategy is required to better assess the nature of these sources. ",
        "introduction": "The X--ray emission from Seyfert host galaxies comprises the contribution of a number of discrete sources plus the hot interstellar plasma (Fabbiano 1989). Most of the discrete sources appear to be close accreting binaries, with a compact companion. \\emph{Einstein} observations of the bulge of M31 revealed a population of about 100 low mass X--ray binaries (Fabbiano et al. 1987). Later, Supper et al. (1997) showed, by using \\emph{ROSAT} data, that the most luminous of these objects in M31 has $L_X=2\\times 10^{38}$ erg s$^{-1}$, close to the Eddington limit for a $1.4 M_{\\sun}$ neutron star.  Recently, several sources with X--ray luminosities higher than the Eddington limit for a typical neutron star have been detected in nearby galaxies (e.g., Read et al. 1997; Colbert \\& Mushotzky 1999; Makishima et al. 2000; La Parola et al. 2001; Zezas et al. 2001). Fabbiano et al. (2001) found with \\emph{Chandra} 14 pointlike sources in the Antennae galaxies, with luminosities above $10^{39}$ erg s$^{-1}$ and up to $10^{40}$ erg s$^{-1}$.  These discoveries have raised difficulties in the interpretation of these sources. Even though it is statistically possible to have \\emph{some} individual cases of off--centre black holes with masses of the order of $10^{3}-10^{4}$ $M_{\\sun}$ (by assuming a typical Eddington ratio of $0.1-0.01$; cf. Nowak 1995), it is very difficult to explain the high number of sources detected so far within this scenario.  Dynamical friction should have caused the objects to spiral to the nucleus of the galaxy. Several other hypotheses have been suggested about the nature of ULX: anisotropic emission from accreting black holes (King et al. 2001), emission from jets in microblazars (K\\\"ording et al. 2002), emission from accreting Kerr black holes (Makishima et al. 2000), and inhomogeneities in radiation--pressure dominated accretion disks (Begelman 2002). However, the lack of sufficient information has not allowed us to distinguish between the different models proposed. The search for optical counterparts has not yet yielded much data: to date only one ULX appears to have a plausible counterpart (Roberts et al. 2001), and other ULX may be associated with planetary nebulae or H~II regions (Pakull \\& Mirioni 2002; Wang 2002).  Our team has been awarded about 250 ks of XMM--EPIC guaranteed time, and we started a distance--limited survey of Seyfert galaxies. We selected $28$ objects in the northern hemisphere with $B_{\\mathrm{T}}<12.5$ mag and $d<22$~Mpc (Di Cocco et al. 2000; Cappi et al. 2002) from the Palomar survey of Ho et al. (1997a). The distances were estimated according to Ho et al. (1997a), and we adopt the same convention in the present paper.  Here we present the results from a study of the discrete sources detected in the galaxies, which are neither the nucleus nor background objects. To date we have obtained 13 objects in our sample, but three observations were heavily corrupted by soft--proton flares and it was not possible to extract any useful information. Here we present part of a study of the discrete source population in the remaining 10 Seyfert galaxies (Table~\\ref{tab:host}), specifically the catalog of ULX sources. Some preliminary results have been presented in Foschini et al. (2002).  Previous detections of ULX have been largely confined to late--type galaxies (e.g., IC 342, M82, NGC 3628, and NGC 5204) or interacting systems undergoing a starburst phase (e.g., the Antennae). Although the objects studied here technically have Seyfert nuclei, the level of nuclear activity is extremely low, and for the present purposes they can be considered ``typical'' nearby galaxies.  The one selection effect to bear in mind is that most of the Palomar Seyferts tend to be relatively bulge--dominated disk galaxies (see, e.g., Ho et al. 1997b), and so late--type galaxies are underrepresented in our sample.  \\begin{table*}[!ht] \\caption{Main characteristics of the observed host galaxies. Columns: (1) Name of the host galaxy from the New General Catalog; (2) optical coordinates of the nucleus from Cotton et al. (2001); (3) Hubble type; (4) spectral classification of the nucleus; (5) distance [Mpc]; (6) major axis of the $D_{25}$ ellipse [arcmin]; (7) Galactic absorption column density [$10^{20}$~cm$^{-2}$]; (8) date of \\emph{XMM--Newton} observation [year--month--day]; (9) effective exposure time, i.e. cleaned from soft--proton flares [ks]. Data for Cols.~$3-6$ are taken from Ho et al. (1997a).} \\centering \\begin{tabular}{lcccccccc} \\hline Galaxy & R.A., Dec. (J2000)  & Hubble Type & Sp. Class. & $d$ & $D_{25}$ & $N_{\\mathrm{H}}$ & Date & Exp.\\\\ (1)    & (2)              & (3)         & (4)        & (5) & (6)      & (7)     & (8)        & (9)\\\\ \\hline NGC1058 & 02:43:30.2, +37:20:27.2     & SA(rs)c     & S2            & 9.1      & 3.02 &  6.65 & $2002-02-01$ & 6.0\\\\ NGC3185 & 10:17:38.7, +21:41:17.2     & SB(r)0/a    & S2        & 21.3     & 2.34 &  2.12 & $2001-05-07$ & 9.1\\\\ NGC3486 & 11:00:24.1, +28:58:31.6     & SAB(r)c     & S2            & 7.4      & 7.08 &  1.9  & $2001-05-09$ & 4.2\\\\ NGC3941 & 11:52:55.4, +36:59:10.5     & SB(s)0      & S2        & 18.9     & 3.47 &  1.9  & $2001-05-09$ & 5.0\\\\ NGC4138 & 12:09:29.9, +43:41:06.0     & SA(r)0+     & S1.9          & 17.0     & 2.57 &  1.36 & $2001-11-26$ & 10.0\\\\ NGC4168 & 12:12:17.3, +13:12:17.9     & E2          & S1.9      & 16.8     & 2.75 &  2.56 & $2001-12-04$ & 17.4\\\\ NGC4501 & 12:31:59.3, +14:25:13.4     & SA(rs)b     & S2            & 16.8     & 6.92 &  2.48 & $2001-12-04$ & 2.8\\\\ NGC4565 & 12:36:21.1, +25:59:13.5     & SA(s)b      & S1.9          & 9.7      & 15.85&  1.3  & $2001-07-01$ & 10.0\\\\ NGC4639 & 12:42:52.5, +13:15:24.1     & SAB(rs)bc   & S1.0          & 16.8     & 2.75 &  2.35 & $2001-12-16$ & 9.7\\\\ NGC4698 & 12:48:23.0, +08:29:14.8     & SA(s)ab     & S1.9          & 16.8     & 3.98 &  1.87 & $2001-12-16$ & 9.2\\\\ \\hline \\end{tabular} \\label{tab:host} \\end{table*}  \\begin{table*}[!ht] \\caption{ULX in the present catalog (significance greater than $4\\sigma$). Columns: (1) Name of the host galaxy; (2) ULX number; (3) coordinates of the ULX; (4) angular separation from the optical centre of the galaxy [arcsec]; (5) ULX name according to \\emph{XMM-Newton} rules.} \\centering \\begin{tabular}{lcccc} \\hline Host Galaxy & Object & RA, Dec (J2000.0)  & Separation & {\\it XMM} ID\\\\ (1)         & (2)    & (3)                & (4)      & (5)\\\\ \\hline NGC1058 & ULX1    & 02:43:23.5, +37:20:38           & 77           & XMMU J024323.5+372038 \\\\ {}      & ULX2    & 02:43:28.3, +37:20:23           & 19           & XMMU J024328.3+372023 \\\\ NGC3185 & ULX1    & 10:17:37.4, +21:41:44           & 30           & XMMU J101737.4+214144 \\\\ NGC3486 & ULX1    & 11:00:22.4, +28:58:18           & 23           & XMMU J110022.4+285818 \\\\ NGC3941 & ULX1    & 11:52:58.3, +36:59:00           & 38           & XMMU J115258.3+365900 \\\\ NGC4168 & ULX1    & 12:12:14.5, +13:12:48           & 45           & XMMU J121214.5+131248 \\\\ NGC4501 & ULX1    & 12:32:00.1, +14:22:28           & 166          & XMMU J123200.1+142228 \\\\ {}      & ULX2    & 12:32:00.8, +14:24:42           & 40           & XMMU J123200.8+142442 \\\\ NGC4565 & ULX1    & 12:36:05.2, +26:02:34           & 289          & XMMU J123605.2+260234 \\\\ {}      & ULX2    & 12:36:14.8, +26:00:53           & 127          & XMMU J123614.8+260053 \\\\ {}      & ULX3    & 12:36:17.3, +25:59:51           & 59           & XMMU J123617.3+255951 \\\\ {}      & ULX4    & 12:36:17.4, +25:58:54           & 51           & XMMU J123617.4+255854 \\\\ {}      & ULX5    & 12:36:18.8, +26:00:34           & 83           & XMMU J123618.8+260034 \\\\ {}      & ULX6    & 12:36:27.8, +25:57:34           & 139          & XMMU J123627.8+255734 \\\\ {}      & ULX7    & 12:36:30.6, +25:56:50           & 197          & XMMU J123630.6+255650 \\\\ NGC4639 & ULX1    & 12:42:48.3, +13:15:41           & 61           & XMMU J124248.3+131541 \\\\ {}      & ULX2    & 12:42:51.4, +13:14:39           & 50           & XMMU J124251.4+131439 \\\\ NGC4698 & ULX1    & 12:48:25.9, +08:30:20           & 73           & XMMU J124825.9+083020 \\\\ \\hline \\end{tabular} \\label{tab:srccoord} \\end{table*}  \\begin{table*}[!ht] \\caption{X--ray data on the ULX of the present catalog. Columns: (1) Host galaxy; (2) ULX number; (3) count rate in counts per second, as calculated with \\emph{eboxdetect}; (4) likelihood of the detection (cf. Ehle et al. 2001); (5) absorbing column density [$10^{21}$ cm$^{-2}$]; (6) photon index of the power law; (7) X--ray luminosity in the $0.5-10$ keV energy band [$10^{38}$ erg s$^{-1}$], calculated from the count rates of Col.~3 and converted with the factor of $3\\times 10^{11}$~cnt$\\cdot$cm$^2$/erg, which in turn is derived by using the power-law model with $\\Gamma=2.0$ and an average Galactic $N_{\\mathrm{H}}=3\\times 10^{20}$ cm$^{-2}$ (cf. Ehle et al. 2001). When the parameters $N_H$ and $\\Gamma$ are present, it means that the statistics are sufficient to perform a spectral fitting and uncertainties in the parameters estimate are at the $90\\%$ confidence limits. For $N_{\\mathrm{H}}=N_{\\mathrm{H, gal}}$, it means that no additional absorption is found to be significant. The luminosities were calculated using distances from Ho et al. (1997a), who considered an infall velocity of 300 km/s for the Local Group, $H_0=75$~km$\\cdot$s$^{-1}$Mpc$^{-1}$, and the distance of the Virgo Cluster of 16.8~Mpc (cf. Table~1, Col.~5).} \\centering \\begin{tabular}{lcccccc} \\hline Host Galaxy & Object & Count Rate $[10^{-3}]$  & Likelihood & $N_{\\mathrm{H}}$  & $\\Gamma$ & $L_{0.5-10 keV}$\\\\ (1)         & (2)    & (3)                              & (4)        & (5)      & (6)      & (7)\\\\ \\hline NGC1058 & ULX1    & $13\\pm 2$                           & 79         & $N_{\\mathrm{H, gal}}$ & $1.1\\pm 0.3$  & $11$\\\\ {}      & ULX2    & $6\\pm 1$                            & 22         & --       & --       & $2.3$\\\\ NGC3185 & ULX1    & $5\\pm 1$                            & 27         & --       & --       & $13$\\\\ NGC3486 & ULX1    & $20\\pm 4$                           & ($^{\\mathrm{*}}$) & $N_{\\mathrm{H, gal}}$ &  $2.2\\pm 0.5$ & $5.0$\\\\ NGC3941 & ULX1    & $43\\pm 4$                           & 320        &  $N_{\\mathrm{H, gal}}$ & $1.9\\pm 0.2$ & $74$\\\\ NGC4168 & ULX1    & $4.0\\pm 0.7$                        & 36         & --       & --       & $6.0$\\\\ NGC4501 & ULX1    & $11\\pm 3$                           & 23         & --       & --       & $17$\\\\ {}      & ULX2    & $29\\pm 4$                           & 93         & $N_{\\mathrm{H, gal}}$ & $2.3\\pm 0.4$  & $37$\\\\ NGC4565 & ULX1    & $7\\pm 1$                            & 44         & --       & --       & $3.4$\\\\ {}      & ULX2    & $10\\pm 1$                           & 90         & $6\\pm 5$ & $1.7\\pm 0.6$ & $16$\\\\ {}      & ULX3    & $3.9\\pm 0.9$                        & 21         & --       & --       & $2.0$\\\\ {}      & ULX4    & $51\\pm 3$                           & 821        & $N_{\\mathrm{H, gal}}$ & $1.9\\pm 0.1$  & $25$\\\\ {}      & ULX5    & $7\\pm 1$                            & 57         & --       & --       & $3.4$\\\\ {}      & ULX6    & $13\\pm 2$                           & 121        & $N_{\\mathrm{H, gal}}$ & $1.5\\pm 0.3$  & $9.0$\\\\ {}      & ULX7    & $4\\pm 1$                            & 23         & --       & --       & $2.0$\\\\ NGC4639 & ULX1    & $5\\pm 1$                            & 29         & --       & --       & $8.0$\\\\ {}      & ULX2    & $3\\pm 1$                            & ($^{\\mathrm{*}}$)  & -- & --     & $5.0$\\\\ NGC4698 & ULX1    & $16\\pm1$                            & 168        & $N_{\\mathrm{H, gal}}$ & $2.0\\pm 0.2$     & $30$\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{*}}$] Data from manual analysis. \\end{list} \\label{tab:srcdata} \\end{table*} ",
        "conclusions": "We have analyzed the X--ray data from \\emph{XMM--Newton} observations of 10 nearby Seyfert galaxies. The host galaxies are located between $7.4$ and $21.3$~Mpc, with 7/10 between $16.8$ and $21.3$~Mpc. Only one host galaxy (NGC4168) is elliptical (E2), while all the remaining are spirals of various types (see Table~\\ref{tab:host}).  We found ULX in 9 of the 10 galaxies. The contamination with background sources is very low. Indeed, from \\emph{XMM--Newton} observations of the Lockman Hole, Hasinger et al. (2001) found about $100$ sources per square degree with flux higher than $10^{-14}$~erg cm$^{-2}$ s$^{-1}$ in the energy band $0.5-2$~keV, and about $200$ sources per square degree in the energy band $2-10$~keV. Assuming the same $\\log N - \\log S$ and considering the flux limit reached by our observations, we expect to find, in the worst case (NGC4565), fewer than one (0.7) background source inside the $D_{25}$ ellipse. For all remaining host galaxies, the expected number of background objects is significantly less than one (0.2 for most of the cases). Therefore, we expect that, in the worst case, the overall sample contains fewer than 2 background objects.  The total number of ULX in the present catalog is 18. The mean value is 1.8 ULX per galaxy, with the maximum value in NGC4565 with 7 ULX and the minimum in NGC4138 with no detection. By omitting NGC4565, we have a mean value of about 1.2 ULX per galaxy. With respect to the {\\it ROSAT}/HRI survey by Roberts \\& Warwick (2000), we find that \\emph{XMM--Newton} allows a significative improvement in the number of ULX detections (see Table~\\ref{tab:inc}). In Table~\\ref{tab:inc} we list also the luminosities in $B$ and far-infrared bands of the host galaxies, the latter being a rough indicator of the star formation activity. At a first look, no evident correlation appears, but the present sample is very small. Although it appears that host galaxies with high $L_{\\mathrm{B}}$ and $L_{\\mathrm{FIR}}$ (NGC4501 and NGC4565) have a higher number of ULX, we caution that this effect could be spurious because our survey does not reach a uniform luminosity threshold to detected ULX.  NGC4565, for example, reaches a flux limit of $\\sim$$1\\times 10^{-14}$~erg cm$^{-2}$ s$^{-1}$, which is deeper than the flux corresponding to the ULX luminosity limit of $2\\times 10^{38}$~erg s$^{-1}$ (cf. Sect.~2). For NGC4501, on the other hand, the flux limit is too shallow to reach this luminosity limit.  \\begin{table}[!ht] \\caption{Predicted and observed number of ULX. Columns: (1) Name of the host galaxy; (2) total absolute $B$ magnitude, from Ho et al. (1997a); (3) luminosity [$10^{10}$~$L_{\\sun}$~erg s$^{-1}$] calculated from data in Col.~2; (4) expected number of ULX according to Roberts \\& Warwick (2000), who found a relationship between the number of ULX and $L_{\\mathrm{B}}$; (5) number of ULX actually found in the present survey; (6) luminosity in the far-infrared (FIR, $42.5-122.5$~$\\mu$m) in units [$10^{42}$~erg s$^{-1}$], calculated from data of Ho et al. (1997a).} \\centering \\begin{tabular}{lccccc} \\hline Galaxy & $M_{\\mathrm{B}}$      & $L_{\\mathrm{B}}$      & Expected   & Found & $L_{\\mathrm{FIR}}$\\\\ (1)    & (2)          & (3)          & (4)        & (5)   & (6)\\\\ \\hline NGC1058 & $-18.25$      & 0.15        & 0.1        & 2    & 1.9\\\\ NGC3185 & $-18.99$      & 0.30        & 0.2        & 1    & 5.1\\\\ NGC3486 & $-18.58$      & 0.21        & 0.1        & 1    & 2.7\\\\ NGC3941 & $-20.13$      & 0.87        & 0.6        & 1    & 5.3$^{\\mathrm{*}}$\\\\ NGC4138 & $-19.05$      & 0.32        & 0.2        & 0    & 2.0$^{\\mathrm{*}}$\\\\ NGC4168 & $-19.07$      & 0.32        & 0.2        & 1    & 0.37\\\\ NGC4501 & $-21.27$      & 2.5         & 1.7        & 2    & 48\\\\ NGC4565 & $-20.83$      & 1.7         & 1.2        & 7    & 10\\\\ NGC4639 & $-19.28$      & 0.40        & 0.3        & 2    & 4.1\\\\ NGC4698 & $-19.98$      & 0.70        & 0.5        & 1    & 1.5\\\\ \\hline \\end{tabular} \\label{tab:inc} \\begin{list}{}{} \\item[$^{\\mathrm{*}}$] Since no FIR data were available in Ho et al. (1997a), we calculate the FIR luminosity according to the relationship $\\log L_{\\mathrm{FIR}}/L_{\\mathrm{B}}=-0.792$ (Pogge \\& Eskridge 1993). \\end{list} \\end{table}   The luminosities observed are in the range $(2 - 74)\\times 10^{38}$ erg s$^{-1}$, depending on the model considered. If we make the simplistic assumptions that the accretion is uniform and spherical, that the bolometric luminosity approximately equals the X--ray luminosity, and that the Eddington ratio is $1$, these luminosities correspond to compact objects with masses between $1.5$ and $57$ $M_{\\sun}$. However, the X--ray luminosity is generally only $30-40\\%$ of the bolometric luminosity of the accreting sources (e.g., Mizuno et al. 1999). In addition, if the Eddington ratio is in the range of $0.1-0.01$, as suggested by observations of Galactic black hole candidates (e.g., Nowak 1995), the mass range would shift toward $10^3-10^4$ $M_{\\sun}$.  Unless the sources are very young, such high masses are difficult to explain for off--centre sources, because dynamical friction would tend to drag the objects toward the centre in less than the Hubble time (cf. Binney \\& Tremaine 1987).  Therefore, as proposed by several authors, alternative scenarios must be considered. For example, Makishima et al. (2000) proposed a Kerr black hole scenario: in this case, the luminosity produced by a spinning black hole can be up to 7 times larger than in a Schwarzschild black hole. On the other hand, King et al. (2001) suggested that the matter could accrete anisotropically: an anisotropic factor of $0.1-0.01$ reduces the values of the mass to those typically observed in X--ray binaries in the Milky Way. Something similar has been suggested by Begelman (2002): in this case, the presence of inhomogeneities in radiation pressure--dominated accretion disks, as a consequence of photon--bubble instability, would allow the radiation to escape. Finally, K\\\"ording et al. (2002) and Georganopoulos et al. (2002) suggested the possibility of relativistic beaming due to the presence of jets coupled to an accretion disk. Both are based on the microquasar model by Mirabel \\& Rodriguez (1999).  The statistics of the present observations do not allow us to discriminate clearly between the different models, but we can infer some useful hints from the eight sources, which gave sufficient counts for a spectral fitting. In 5/8 cases the best-fit model is obtained with a simple power law with $\\Gamma \\approx 1.9-2.3$ (for an example of spectra, see Fig.~\\ref{fig4}). One of these five sources (NGC1058--ULX1) presents an almost flat spectrum ($\\Gamma = 1.1\\pm 0.3$). For the remaining sources (2/8), we obtained a best fit with the black body model with $kT\\approx 0.5-0.9$~keV.  It is known that the emission expected from a black hole X--ray binary is variable: in the hard state, the spectrum is typically a power law with $\\Gamma \\approx 1.3-1.9$, while in the soft state the spectral index increases up to about 2.5 and a soft component appears in the X--ray spectrum (e.g., Ebisawa et al. 1996). Therefore, our sources could be black hole X--ray binaries in a hard or soft state. Terashima \\& Wilson (2002) proposed the existence of two different populations of ULX, one characterized by soft thermal and the other by non--thermal X--ray emission. A possible key to distinguish between the available hypotheses can be to perform time variability studies, but current statistics are too low for such a study.  It is interesting to note that the MCD model, which has often been successful in the past for ULX (e.g., Colbert \\& Mushotzky 1999; Makishima et al. 2000) is never the best fit in our data. Even when we obtain a reasonable fit with MCD, a simple black body model is statistically better. This may be due to the low photon counts of the present spectra.  The unsaturated Comptonization (CST) model does not give acceptable fits, except for ULX4 in NGC4565, for which it represents the second best fit, after the power law.  Our understanding of the nature of ULX is limited by the fact that, to date, counterparts at other wavelengths are quite rare (Roberts et al. 2001; Pakull \\& Mirioni 2002; Wang 2002). In our sample, we note that one source (NGC4698--ULX1) is detected in the radio (6 cm). NGC4168--ULX1 and NGC4639--ULX2 both show a highly probable optical counterpart. In the second case, it is identified as a H~II region, which is a type of counterpart frequently associated with ULX (Pakull \\& Mirioni 2002). It is worth noting the case of NGC4656--ULX3 could be obscured by, rather than correlated with, a dust cloud. Finally, the probable counterparts of NGC4168--ULX1 and NGC4698--ULX1 appear to be considerably red objects, with $B-R > 2$ mag."
    },
    "0206/astro-ph0206242_arXiv.txt": {
        "abstract": "{We present 3D, gasdynamic simulations of jet/cloud collisions, with the purpose of modelling the HH~270/110 system. From the models, we obtain predictions of H$\\alpha$ and H$_2$ 1-0 s(1) emission line maps, which qualitatively reproduce some of the main features of the corresponding observations of HH~110. We find that the model that better reproduces the observed structures corresponds to a jet that was deflected at the surface of the cloud $\\sim 1000$~yr ago, but is now boring a tunnel directly into the cloud. This model removes the apparent contradiction between the jet/cloud collision model and the lack of detection of molecular emission in the crossing region of the HH~270 and HH~110 axes. ",
        "introduction": "The Herbig-Haro (HH) jet HH~110 (discovered by Reipurth \\& Bally 1986) is the best observed example of a possible HH jet/dense cloud collision. Reipurth et al. (1996) have interpreted the rather unique, collimated but quite chaotic structure of HH~110 as the result of a deflection of the faint HH~270 jet through a collision with a dense cloud.  \\begin{figure*} \\centering \\includegraphics{MS2563f1.eps} \\caption{Time sequence of the density stratifications obtained from model A. The density stratifications on the $y=0$ plane (which includes the outflow axis and the centre of the spherical cloud) are shown for different integration times (as indicated at the bottom left of each plot). The densities are depicted with a logarithmic greyscale, with the values given (in g~cm$^{-3}$) by the bar on the top left of the figure. The $x$ (horizontal) and $z$ (vertical) axes are labeled in cm.} \\label{dens} \\end{figure*}  The evidence presented by Reipurth et al. (1996) for this interpretation can be summarized as follows: \\begin{itemize} \\item no stellar source has been detected aligned with the HH~110 jet, \\item the HH~270 jet (ejected from a detected IR and radio source, see Rodr\\'\\i guez et al. 1998) points towards the beginning of the HH~110 jet, \\item the proper motions of HH~270 and HH~110 have an approximately 2 to 1 ratio, which is completely consistent with the $\\approx 60^\\circ$ deflection angle defined by the locci of the two jets (the flow approximately lying on the plane of the sky). \\end{itemize} This last statement can be understood as follows. When a radiative jet hits the surface of a dense cloud at an incidence angle $\\phi$ (between the incident jet axis and the cloud surface), the normal component of the jet velocity is stopped in a radiative shock, and the jet continues to flow parallel to the surface (conserving the component of the incident jet velocity parallel to the surface). Therefore, the velocity $v_{def}$ of the deflected jet is approximately equal to the projection of the incident jet velocity $v_{inc}$ along the surface of the dense cloud. Therefore, $v_{def}\\approx v_{inc}\\cos\\phi$. One can straightforwardly see that the proper motions and deflection angle defined by HH~270 and HH~110 (see above) do satisfy this condition.  This result led Raga \\& Cant\\'o (1995) to study the dynamics of the collision of a radiative, HH jet with the surface of a dense cloud. These authors presented an analytic model and plane, 2D simulations of the early stages of such an interaction, and found that the general characteristics of the HH~270/110 system could be explained in terms of such a model. The main problem found with the models is that in a rather short timescale, the incident jet starts to perforate the obstacle, and the deflected jet beam is then pinched off. In order to obtain a long enough timescale for the production of the deflected jet, it is necessary to have a very high cloud-to-jet density ratio. Raga \\& Cant\\'o (1995) suggested that this problem might be overcome if the incident jet did not have a completely straight jet beam, so that the impact point would roam over the surface of the dense cloud.  \\begin{figure} \\centering \\includegraphics[width=8cm]{MS2563f2.eps} \\caption{Temperature stratification (top) and adaptive grid structure (bottom) on the $y=0$ plane obtained from model A for a $t=2600$~yr integration time. The temperature stratification is depicted with a logarithmic greyscale with the values given (in K) by the bar on the top of the figure. In the bottom plot, two thick lines separating the jet, cloud and environmental material are shown (see the text). These lines show two values of the passive scalar~: $\\psi=0$ (outer contour) and $\\psi=1.5$ (inner contour). } \\label{temp} \\end{figure}  The regime in which the jet has punched a hole through a cloud was described by Cant\\'o \\& Raga (1996) and Raga \\& Cant\\'o (1996). If the cloud is stratified, the path of the jet through the cloud is curved, though the curvature is important only if the radius of the cloud is comparable to the jet radius. 3D gasdynamic simulations of the penetration of a jet into and through a dense, stratified cloud were carried out by de Gouveia Dal Pino (1999).  Finally, Hurka et al. (1999) have studied the bending of the beam of a 3D MHD, non-radiative jet by a magnetic field with a strong gradient (as would be found at the surface of a dense cloud). These authors show that this effect would help to increase the timescale over which the jet/cloud surface interaction takes place, before the deflected jet is pinched off.  In the present paper, we discuss 3D gasdynamic simulations of the interaction of a radiative jet with the surface of a dense cloud. We show the results from two simulations with different assumptions for the incident jet~: \\begin{itemize} \\item that the jet is ejected with a constant direction and velocity \\item that it is produced with a precessing outflow direction and a sinusoidally varying velocity. \\end{itemize} Through a comparison of these two simulations, we can evaluate the effect of a ``roving'' impact point on the production of the deflected jet.  \\begin{figure} \\centering \\includegraphics[width=8cm]{MS2563f3a.eps} \\includegraphics[width=8cm]{MS2563f3b.eps} \\caption{Constant density surface (corresponding to a $n=20$~cm$^{-3}$ number density of atomic nuclei) from model A for a $t=2600$~yr integration time. The two graphs show the surface as seen from two different directions.} \\label{3d} \\end{figure}   Our simulations include a treatment of the dissociation and ionization of the gas. Therefore, we can use the results to obtain predictions of atomic and molecular lines, which we directly compare with previously published images of HH~110.  In particular, we compute predicted maps in the H$_2$ 1-0 s(1) line. This is of interest because the morphology of HH~110 in this IR line is quite strikingly different from its morphology in atomic/ionic lines. Davis et al. (1994) and Noriega-Crespo et al. (1996) found that the H$_2$ emission is much better collimated, and lies along one of the edges of the HH~110 jet beam. This led Noriega-Crespo et al. (1996) to present a simple model of the molecular emission as coming from material from the dense cloud which is entrained by the jet as it brushes past the cloud surface. Our present simulations allow us to make a more definite assesment of whether or not such a mechanism actually succeeds in explaining the molecular emission of HH~110.  We should point out that Choi (2001) present HCO$^+$ emission maps, in which they detect emission in the HH~270/110 region, but not around the ``point of impact'' in which the ``incident'' HH~270 jet is presumably redirected into the ``deflected'' HH~110 jet. This result leads them to suggest that HH~110 might actually not be the result of a jet/cloud collision, but that it could instead be a ``straight'' jet ejected from a low luminosity, undetected stellar source which is presumably more or less aligned with the direction of the HH~110 flow. In the conlcusions, we discuss the possible ways of reconciling the jet/cloud interaction model with the observations of Choi (2001) which are suggested by our 3D gasdynamic simulations. ",
        "conclusions": "We have presented two jet/cloud collision 3D gasdynamic simulations~: one with an incident jet with time-independent injection conditions (model A), and a second one with a variable velocity, precessing incident jet (model B). A $\\rho_c/\\rho_j=100$ cloud to initial jet density ratio has been chosen for both models.  Model A produces a deflection of the jet beam only for a $\\sim 500$-1000~yr timescale, after which the incident jet starts digging a straight tunnel through the dense cloud. At later times, the deflected jet material continues to travel away from the impact region, leaving behind a complex ``wake''.  Model B produces a broader jet/cloud impact region as a result of the jet precession. This effect results in a longer timescale for the duration of the jet deflection on the cloud surface (in fact, the jet deflection is still occuring at the end of our $t=3000$~yr numerical simulation).  Model B is more successful at reproducing the proper motions of HH~110, giving $\\sim 100$~km~s$^{-1}$ velocities for the H$\\alpha$ intensity maxima along the deflected jet beam. Model A gives velocities of $\\sim 15$-45~km~s$^{-1}$, which are substantially lower than the proper motion velocities of HH~110 (see Reipurth et al. 1996).  This difference between model A and model B is due to two effects. The first effect is that in model A, the jet is deflected only for a $\\sim 500$-1000~yr timescale, and that this deflected jet material then slows down as it interacts with surrounding, environmental gas. In model B, successive deflected ``bullets'' (i.~e., internal working surfaces) travel into the low density region left behind by the passage of the head of the deflected jet, and do not interact directly with the higher density environment. The second effect is that because of the precession of model B, some of the bullets have trajectories which are more tangential to the surface of the molecular cloud. These more tangential bullets are less deflected, and therefore preserve larger velocities than the ones that have a more normal incidence on the cloud surface (or than the deflected jet of model A).  However, in most other counts, model A is more successful than model B at reproducing the observations of HH~110~: \\begin{itemize} \\item the general qualitative appearance of the H$\\alpha$ maps is in better agreement with the HH~110 H$\\alpha$ images, \\item the features of the H$_2$ 1-0 s(1) emission and their spatial relation to the H$\\alpha$ emission also are in good qualitative agreement with HH~110, \\item the maps that better resemble HH~110 correspond to times (e.~g., the $t=2400$~yr frame of figures 1 and 6) in which the impact region is already immersed within the cloud. \\end{itemize} This last feature offers an interesting way of reconciling the jet/cloud collision model with the HCO$^+$ observations of Choi (2001).  In these observations, no HCO$^+$ emission was detected in the region in which the axis of the ``incident'' HH~270 jet crosses the axis of the ``deflected'' HH~110. Choi (2001) noted that this appeared to be in disagreement with a jet/cloud collision model for this system, as the cloud shock produced in the impact region should indeed produce HCO$^+$ emission.  The situation found in model A, however, could indeed be in agreement with the observations of Choi (2001). In this model, the impact region does not lie in the point in which the incident and deflected jet axes cross, but is instead located further along the axis of the incident jet. Interestingly, Choi (2001) does find substantial HCO$^+$ emission West of HH~110, approximately aligned with HH~270.  As we have discussed in \\S 3.2, the fact that the impact region is not observed optically can in principle be a result of the dust extinction in the dense cloud. Interestingly, some H$_2$ 1-0 s(1) emission is apparently detected to the West of HH~110 (more or less aligned with HH~270, see Noriega-Crespo et al. 1996), which in principle might be associated with the impact region.  To conclude, we note the two main features of our results~: \\begin{itemize} \\item we find that the jet/cloud interaction model does reproduce the H$\\alpha$ and H$_2$ 1-0 s(1) emission observed in HH~110 in a qualitatively successful way, \\item our models show that the lack of a detected impact zone in the incident/deflected jet crossing regions (Choi 2001) is not a major problem for a jet/cloud collision model (as this region could presently be displaced further into the cloud). \\end{itemize}  Clearly, important questions remain about the more technical aspects of our simulations. In particular, the H$_2$ emission from the deflected jet comes from molecular cloud material which has been entrained into the jet flow (see figures 2 and 5). Even though we find that our models produce H$_2$ emission structures in agreement with the observations of HH~110, the accuracy of our rather low resolution simulations in reproducing the entrainment process (which gives rise to this emission) is somewhat questionable. Because of this, the present results have to be taken with some caution."
    },
    "0206/astro-ph0206074_arXiv.txt": {
        "abstract": "A search for faint slowly variable objects was undertaken in the hope of finding QSO candidates behind the Small and Large Magellanic Clouds (SMC and LMC). This search used the optical variability properties of point sources from the Magellanic cloud OGLE-II photometric data. Objects bluer than $V-I=0.9$ and within $ 17 < I < 20.5 $ were studied. Robust variograms/structure functions have been computed for each time-series and only candidates showing a significant increasing variability over longer time scales were selected.  Several light curves were identified as having probable artifacts and were therefore removed.  Stars showing signs of periodicity or small trends in their light curves were also removed and we are left with mostly either Be stars ($\\gamma\\,$Cas stars) or QSO candidates.  We present a list of 25 slowly varying objects for SMC and 155 for LMC, out of 15'000 and 53'000 variable objects respectively. Of these, about 15 objects for the SMC and 118 objects for the LMC are QSO candidates. ",
        "introduction": "The main motivation of this study is to find QSOs behind the Magellanic clouds. Few QSOs behind SMC and LMC are known.  Several previous studies have reported such discoveries with their associated coordinates, but none of them was in OGLE fields. However more recently, Dobrzycki et al. (2002) found four new QSOs thanks to spectroscopic observations of the optical counterpart of X-ray sources. Those four candidates were observed by OGLE-II.  Other studies have been undertaken for several years by the MACHO group (Geha et al. 1999, Drake et al. 2001). They found about 30 QSOs behind LMC, but have not yet published their coordinates.   QSOs can be useful for different purposes, for example: for fixing the coordinate system in proper motion studies of LMC, SMC or foreground stars; for mapping the interstellar material in the Magellanic Clouds.  The optical variability properties of QSOs are not very well known. Some efforts are underway to remedy this situation; for instance AGNs are monitored in 11 passbands by Kobayashi et al. 1998 (MAGNUM project).  Most QSOs are variable (Cristiani et al. 1996). Discovery of QSOs have been achieved using their optical variability (Trevese et al. 1989, Brunzendorf \\& Meusinger 2001).  We note however that only one QSO out of the four in Dobrzycki's studies was classified as a variable in OGLE-II. This is mainly due to the faintness of the three other objects (about $V \\simeq 20$) and the resulting large errors ($\\simeq 0.1$) in the OGLE-II photometry.  We exploit the variability properties of QSOs since photometric data of OGLE-II is in public domain and offers a 4 years observing period with about 300-400 data points per object in $I$-band and about 30 in $B$ and $V$-bands.  From SDSS, we know (Strauss 2002) that there are about 5 QSOs per square degree brighter than $i = 19$ mag. The number is smaller by a factor 5 per magnitude. As SDSS $i$-mag is approximately $I$-mag, a cut $I > 17$ will exclude one or two QSOs in all fields together, and therefore seems a sensible limit. We should expect to find about 30 QSOs behind the SMC and LMC brighter than $I = 19$ mag. ",
        "conclusions": "The main result of this study shows that it is possible to establish a rather narrow list of QSO candidates using mostly photometric times series in the optical wavelengths. However confirmation with spectroscopic data is needed. We list 118 and 15 QSO candidates for the LMC and SMC respectively.  Once the QSOs are eventually identified, we will be able to fine-tune the algorithm to select QSOs more efficiently.   OGLE-II may give hints of how a sampling could be programmed in order to optimize the detection of QSOs from a variability point of view.  From the study of Dobrzycki et al. (2002), we remark that to select QSOs is not a trivial task. From about one hundred candidates, the 30 best objects were selected for spectroscopic follow-up and 4 objects were confirmed as QSOs.  In a mission like GAIA (Perryman et al. 2001), the number of measured objects is estimated to be of the order of one billion, the distinction between QSOs and stars is subject of study with the current photometric system and astrometric precision (Mignard 2002). Variability could be used as an additional criterion for selecting QSO candidates and therefore diminishing further the rate of false detection.  This study is one additional example, that multiepoch surveys like OGLE, MACHO or EROS originally oriented to detect microlensing events can be used in many different fields of astronomy. The scientific outcomes of such surveys are often unexpected."
    },
    "0206/hep-ph0206157_arXiv.txt": {
        "abstract": "We provide description of rapidity spectra of particles produced in $p\\bar{p}$ collisions using anomalous diffusion approach to account for their non-equilibrium character. In particular, we exhibit connection between multiproduction processes and anomalous diffusion described through the nonlinear Focker-Planck equation with nonlinearity given by the nonextensivity parameter $q$ describing the underlying Tsallis $q$-statistics and demonstrate how it leads to the Feynman scaling violation in these collisions. The $q$ parameter obtained this way turns out to be closely connected to parameter $1/k$ converting the original poissonian multiplicity distribution to its observed Negative Binomial form. The inelasticity of reaction has been also calculated and found to slightly decrease with the increasing energy of reaction. ",
        "introduction": "In description of multiparticle production processes one often uses statistical methods and concepts which follow the classical Boltzmann-Gibbs (BG) approach. However, it was demonstrated recently that to account for the long range correlations and for some intrinsic fluctuations in the hadronizing system one should rather use the nonextensive Tsallis statistics \\cite{T}, in which one new parameter $q$ describes summarily the possible departure from the usual BG case (which is recovered in the $q\\rightarrow 1$ limit) \\cite{WWq}. Here we shall provide detail description of the rapidity spectra of particles (mostly pions) produced in $p\\bar{p}$ collisions using the anomalous diffusion approach to this problem and in this way accounting for their non-equilibrium character.  This method originates from diffusion model approach to nuclear multiparticle production collisions developed in \\cite{Wol}, which has been further (successfully) applied to the recent RHIC data in \\cite{Biya}. The first attempt to extend it also to the case of anomalous diffusion (which corresponds to nonextensive $q\\neq 1$ case) has been presented recently for nuclear collisions \\cite{Lav}. The non-linear Fokker-Planck (FP) equation used in this case has form ($f=f(y,t)$ with $y$ being rapidity and $t$ time variable): \\begin{equation} \\frac{\\delta}{\\delta t}f^{\\mu} = \\frac{\\delta}{\\delta y}\\left\\{ J(y)f^{\\mu} + D \\frac{\\delta}{\\delta y}f^{\\nu}\\right\\}, \\end{equation} $D$ and $J$ are diffusion and drift coefficients respectively. The hadronization process is visualized here as diffusion in the rapidity space starting with rapidities (in cms frame) $Y^{(\\pm)}_{max} \\simeq \\pm \\ln\\frac{\\sqrt{s}}{m_T}$ (where $\\sqrt{s}$ is invariant energy of reaction and $m_T=\\sqrt{m^2 + \\langle p_T\\rangle^2}$ mean transverse mass kept here as given. In the approaches presented in \\cite{Wol} and \\cite{Biya} linear or constant drift coefficients has been used and in effect obtained double gaussian-like form of rapidity spectra\\footnote{Interesingly enough, such spectra for $\\frac{dN}{dy}$ were already postulated and used on purely phenomenological grounds as simple parametrizations of results of string models allowing for fast numerical calculations of cosmic ray cascades \\cite{KlarCap}.}. As was demonstrated in \\cite{Buk,Lav} for drift proportional to the longitudinal momentum of the particle (i.e., $J(y) \\sim \\sinh y$) one gets the thermal (Boltzmann) distribution (for linear FP equation, i.e., for $q=\\nu=1$) whereas for nonlinear-FP equation the time dependent solution has the specific power-like form the norm of which is conserved only for $\\mu=1$, therefore $\\nu=2-q$ and our distribution is given by the following formula: \\begin{equation} f_{q}(y) = \\left[ 1-(1-q)\\frac{m_{T}}{T} \\cosh\\left( y-y_{m} \\right) \\right]^{\\frac{1}{1-q}}. \\label{eq:form} \\end{equation} ",
        "conclusions": "We have provided here description of rapidity spectra of particles produced at CERN and Fermilab energies treating their formation as diffusion process in rapidity space. To account for the anomalous character of such diffusion, the nonlinear form of Focker-Planck equation has been used with nonlinearity described by parameter $q$, the same as the nonextensivity parameter describing the underlying Tsallis statistics. As seen in Figs. 1 and 2, very good agreement with data has been obtained with apparently three parameters: the \"temperature\" $T$, position of the peak at rapidity $y_m$ and parameter $q$ - all logaritmically dependent on the energy of reaction $\\sqrt{s}$ \\footnote{We have checked that we could not get such good fits within formalism with $q=1$, what means that nonlinearity expressed by $q>1$ here is essential feature of the data.}. However, after closer inspection it turns out that parameter $q$, which according to \\cite{WWq} can be regarded as a measure of fluctuations existing in the physical system under consideration, follows essentiall the fluctuations of multiplicity of particles produced at given energy. The small differences noticed in Fig. 3 are, in our opinion, caused entirely by the fluctuations in the inelasticity of the reaction, which makes the initial energy available for the production of secondaries a fluctuating quantity\\footnote{However, we did not attempt here to estimate such fluctuations. We are planning to do it elesewhere.}. Similarly, parameter $y_m$ seems to be closely connected with the maximal available rapidity $Y_{max}(\\sqrt{s})$ \\footnote{There is therefore possibility that it is, in fact, connected with $Y_{max}(K\\cdot \\sqrt{s})$, i.e., it depends in the indirect way on the inelasticty $K$ of the reaction. We have not pursued this problem here.}. This shows that the only parameter which is entirely \"free\" is the limiting temperature $T$ \\footnote{One should notice at this point that its energy dependence as given in eq.(\\ref{eq:param}) contains in itself also the possible energy dependence of the mean transverse mass $m_T$, not accounted for in the formula (\\ref{eq:Fq}).}. Therefore we can say that what we are proposing here is essentially one-parameter fit successfully describing data on rapidity distributions.  The resultant inelasticity $K$ (defined by the formula (\\ref{eq:inel})) turns out to be decreasing function of energy, as seen in Fig. 4, confirming our previous findings in that matter \\cite{inel}. We would like to bring ones attention to the fact that in our approach we can fit equally well (using the same set of parameters) data obtained by P238 \\cite{P238} and UA7 \\cite{UA7} collaborations which cover different regions in rapidity. This makes extrapolation of our formula to higher rapidity region more credible and therefore puts more weight on the obtained energy behaviour of inelasticity.  Acknowledgements: The partial support of Polish Committee for Scientific Research (grants 2P03B 011 18 and 621/E-78/SPUB/CERN/P-03/DZ4/99) is acknowledged.\\\\"
    },
    "0206/astro-ph0206238_arXiv.txt": {
        "abstract": " ",
        "introduction": "Nowadays, the wealth of \\hi\\ data gathered for the Virgo cluster region makes it possible to obtain a reliable description of the pattern of neutral gas deficiency on supercluster scales. I will present here a recent evaluation of the large-scale radial run of the \\hi\\ deficiency on the Virgo~I cluster (VIC) region traced by the giant spiral population, as well as of the distribution of this property in the two- and three-dimensional space, obtained from the combination of 21-cm data with Tully-Fisher (TF) distance measurements. I will also attempt to provide suggestive evidence that objects with high \\hi\\ deficiencies are not exclusively confined to the Virgo cluster proper, but can be also observed both in a background galaxy group at \\sm 25--30 Mpc from us (possible related to the classical W' cloud) and in various galaxies lying in the frontside of the cluster at line-of-sight (LOS) radial distances less than 15 Mpc.  A dynamical model for the collapse and rebound of spherical shells under the point mass and radial flow approximations will be used to demonstrate that it is not unfeasible that some galaxies far from the cluster, including those in the gas-deficient group well to its background, went through the cluster core a few Gyr ago. The implications would be: (1) that a substantial fraction of the \\hi-deficient spirals in the VIC region might have been deprived of their neutral hydrogen by interactions with the hot intracluster medium; and (2) that objects spending a long time outside the cluster cores might keep the gas deficient status without significantly altering their morphology.  \\vspace*{+2mm} ",
        "conclusions": "\\begin{itemize}  \\item Further progress in the knowledge of the detailed structure of the Virgo cluster needs a careful revision of TF distances ---at least until Cepheid distance measurements in Virgo galaxies become more commonplace. Even after the elimination of systematic differences among published Virgo catalogs, a few galaxies still exhibit strongly inconsistent distance measurements: 16 of the 161 members of the 21-cm sample have $1\\sigma$ uncertainties larger than 5 Mpc. The fact that the quoted error is small does not guarantee that the measurement is reliable. If two results clearly disagree, any average, weighted or not, is meaningless, and there is little point in performing it.  \\item There is now compelling evidence of the decisive participation of ram-pressure stripping, which requires high IGM densities and relative velocities, in the reduced gas abundances of the spirals observed in the centers of various rich clusters, either from observational articles \\cite{GJ87,DG91,Sol01} ---including the discovery of galaxies with shrunken \\hi\\ disks \\cite{Cay94,Bra00}--- or from theoretical studies \\cite{SS92,FN99,SAP99,QMB00,Vol01}. The finding that a number of spirals with substantial \\hi\\ deficiencies lie at large radial distances from the Virgo cluster center ---some are very likely members of a background subclump well behind the cluster core--- may seem from the outset hard to reconcile with the proposition that this environmental process is also the cause of their gas deficiencies. However, the modeling of the velocity field around the VIC region demonstrates that characteristics such as a large virgocentric distance or a near turnaround position are not by themselves conclusive indications of a recent arrival. The substantial \\hi\\ deficiency of the background subclump found in \\scite{Sol02} may well have originated on an earlier passage of this entity through the Virgo core. A more precise knowledge of the velocity field around the VIC appears to be required to confirm or discount this possibility.  \\item Even if the tentative suggestion that the \\hi-deficient group on the backside of the VIC might not be a recent arrival is finally proven well-founded, it is still necessary finding a sensible explanation for the apparently long time (\\sm 4--5 Gyr) its gas-poor members have maintained a substantial dearth of gas during without noticeable consequences on their morphologies: 9 of its 15 probable members are late-type spirals, whereas the 5 galaxies with the largest gaseous deficiencies have types Sb or later. Certainly, the details and chronology of the evolution of galactic properties triggered by the sweeping of the atomic hydrogen, as well as its repercussions on the star formation rate, are poorly understood.  \\end{itemize}  The unconspicuous positions of the \\hi-deficient galaxies observed at large virgocentric distances ---the gas-deficient group in the background, for instance, lies halfway between the much larger galaxy concentrations of the Virgo cluster and the background W cloud--- have not encouraged systematic 21-cm surveys outside the central VIC region. I hope that this contribution provides enough grounds for changing the situation and putting the gas content of the galaxies on the outskirts of the Virgo cluster under close examination."
    },
    "0206/astro-ph0206181_arXiv.txt": {
        "abstract": "We analyze the high-resolution X-ray spectrum of Hercules X-1, an intermediate-mass X-ray binary, which was observed with the {\\it XMM-Newton} Reflection Grating Spectrometer. We measure the elemental abundance ratios by use of spectral models, and we detect material processed through the CNO-cycle.  The CNO abundances, and in particular the ratio N/O~$> 4.0$ times solar, provide stringent constraints on the evolution of the binary system.  The low and short-on flux states of Her X-1 exhibit narrow line emission from \\ion{C}{6}, \\ion{N}{6}, \\ion{N}{7}, \\ion{O}{7}, \\ion{O}{8}, \\ion{Ne}{9}, and \\ion{Ne}{10} ions. The spectra show signatures of photoionization. We measure the electron temperature, quantify photoexcitation in the He$\\alpha$ lines, and set limits on the location and density of the gas. The recombination lines may originate in the accretion disk atmosphere and corona, or on the X-ray illuminated face of the mass donor (HZ~Her). The spectral variation over the course of the 35~d period provides additional evidence for the precession of the disk. During the main-on state, the narrow line emission is absent, but we detect excesses of emission at $\\sim$10--15~\\AA, and also near the \\ion{O}{7} intercombination line wavelength. ",
        "introduction": "Hercules X-1 is a bright intermediate-mass X-ray binary which has been observed extensively after its discovery by \\citet[]{first_disc}. The system contains an X-ray pulsar with $P_{\\rm pulse} =1.24$~s period \\cite[]{discovery_pulsar,batse}. A synchrotron resonance feature yields a magnetic field of $B = 3.5 \\times 10^{12}$~G \\cite[]{beppo_cyc,truemper_cyc}. Some models yield a lower magnetic field \\cite[]{lowfield}. Optical light-curves \\cite[]{optical_curves} and X-ray eclipses \\cite[]{discovery_pulsar} yield a $P_{\\rm orb} = 1.7$~day orbital period. The $1.5 \\pm 0.3~M_{\\sun}$ neutron star has a $2.3 \\pm 0.3~M_{\\odot}$ companion, HZ~Her, which changes from A to B spectral-type over the orbital period, due to the strong X-ray illumination on its surface \\cite[]{reynolds}. Pre-eclipse and anomalous dips in the X-ray flux are observed \\cite[]{dipsfound,xteobs} due to the interaction of an accretion stream with the accretion disk. The unabsorbed luminosity of Her X-1 is $L=3.8 \\times 10^{37}$~ergs$^{-1}$, using a distance of $D = 6.6 \\pm 0.4$~kpc \\cite[]{reynolds}.  Hercules X-1 exhibits an unusual long-term X-ray flux modulation with $P_{\\Psi} = 35$~d period \\cite[]{dipsfound}. Changes associated with this period have also been observed in the optical light curves \\cite[]{optical_curves}, X-ray pulse shapes \\cite[]{gingapulse}, X-ray dips \\cite[]{dipsfound,xteobs}, and X-ray spectra \\cite[]{ramsay}. The 35~d X-ray light-curve is asymmetric and contains two maxima: a state of $\\sim 8$~d duration reaching the peak flux $F_{\\rm max}$ named the {\\it main-on}, and a secondary high state of $\\sim 4$~d duration reaching $\\sim 1/3~F_{\\rm max}$ named the {\\it short-on}. A low-flux state with $\\sim 1/20~F_{\\rm max}$ ensues at other epochs. The period $P_{\\Psi}$ varies from cycle to cycle, and it has been observed to be 19.5, 20, 20.5, and 21 times $P_{\\rm orb}$ in the course of $\\gtrsim 5$~yr of continuous monitoring with the All-Sky Monitor onboard the {\\it Rossi X-ray Timing Explorer} \\cite[]{xteobs}. The $\\Psi = 0$ phase is defined as the time the main-on state begins. Observations indicate that $\\Psi =0$ coincides only with orbital phases $\\phi = 0.23$ or 0.68 \\cite[]{xteobs}.  The 35~d cycle has been associated with a tilted accretion disk that precesses by some unknown mechanism \\cite[]{precessing_disk_hyp,twisted_disk_flaps,disk_model}. The X-ray light-curve, the variations in the pulse profiles, and the variability of the dips with $\\Psi$-phase are fit by a geometric model of a precessing, warped accretion disk, together with the beam of a pulsar \\cite[]{disk_model}. Among the parameters obtained in this fit are an $85^{\\circ}$ inclination with respect to the line of sight, a $20^{\\circ}$ precession opening angle for the outermost disk, and an $11^{\\circ}$ precession angle for the innermost disk. The latter precession angles are dependent on assumed input parameters such as the disk thickness.  The ultraviolet (UV) spectrum exhibits line emission from species such as \\ion{C}{4}, \\ion{N}{5}, and \\ion{O}{5}, which originate in two separate components that produce superimposed broad and narrow lines.  The intensity of the narrow lines has the same orbital variation as the UV continuum, which is thought to originate on the illuminated face of HZ~Her. The broad line component follows roughly the velocities expected for the accretion disk \\cite[]{hut_uvlines}. Observations during eclipse ingress and egress suggest that the broad line region originates in a prograde disk of $\\sim 10^{11}$ cm radius \\cite[]{uv_disk_lines}. The UV lines (i.e. \\ion{N}{5}) have a weak P Cygni profile component which indicates the presence of a wind outflow \\cite[]{uv_lines_wind_model}.  The Her X-1 broadband X-ray emission can be described as a blackbody component with temperature $kT \\sim 90$~eV,  plus a power-law component with a 24~keV exponential cutoff, and a 42~keV cyclotron feature \\cite[]{beppo_cyc}. An Fe K fluorescence line evolves with $\\Psi$-phase. The {\\it XMM-Newton} EPIC data, presented by \\cite{ramsay}, show a 6.4~keV Fe K line which is practically unresolved during the low and short-on states, and a broad line at 6.5~keV with $330 \\pm 20$~eV FWHM during the main-on.  In this article, we present the high-resolution spectrum of Her X-1 in the 5 to 38 \\AA \\ band, which we obtained with the \\it XMM-Newton \\rm Reflection Grating Spectrometer (RGS). The high-resolution X-ray spectrum evolves dramatically with 35~d phase.  We analyze the spectral data from three distinct epochs (section \\ref{sec:obs}). The spectrum during the low and short-on states is a faint power-law continuum plus many narrow emission line features from a photoionized plasma (section \\ref{sec:lowshort}).  In contrast, the main-on state has a bright continuum, neutral absorption features, excess of emission near \\ion{O}{7} He$\\alpha$, and a continuum which cannot be fit well with blackbody and/or power-law models (section \\ref{sec:mainon}). The observed emission lines arise from the recombination of electrons with ions and from the subsequent radiative cascades. The gas is predominantly heated by photons due to the large luminosity and small volume of the system.  We analyze and model the high-resolution spectra of Her X-1 to obtain the elemental abundance ratios and measure the state variables of the gas, and we make preliminary identifications of the newly detected emission region(s).  We apply spectral diagnostics on the plasma to constrain its location (section \\ref{sub:helike}), its density (section \\ref{sub:param}), and to measure the electron temperature (section \\ref{sub:rrc}). We introduce a method to extract the elemental abundance ratios from the low and short-on spectra. The method involves extracting the emission measure distribution from model fits of the recombination line fluxes (section \\ref{sec:em}). We discuss whether the narrow line emission originates in a disk atmosphere and corona, on the illuminated companion, or in both (section \\ref{sub:em}).  We describe the implications of the detection of CNO-processed material for the evolution of the X-ray binary (section \\ref{sub:evol}). We outline a physical scenario which produces broad emission lines, which is associated with the magnetopause and the inner disk in X-ray pulsars (section \\ref{sub:magn}). In that context, we speculate on the origin of the complex continuum observed during the main-on (section \\ref{sub:bump}). Our conclusions are in section \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We analyzed the high-resolution X-ray spectra of Her X-1 obtained with the {\\it XMM-Newton} RGS from three observations performed at distinct states of its 35~d cycle.  We detect narrow recombination emission lines during the low and short-on states of Her X-1.  Emission lines are detected from \\ion{C}{6}, \\ion{N}{6}, \\ion{N}{7}, \\ion{O}{7}, \\ion{O}{8}, \\ion{Ne}{9}, and perhaps \\ion{Ne}{10}, plus weak RRC of \\ion{O}{7} and \\ion{N}{7}, which indicate the gas is photoionized. The velocity broadening of the lines is at the resolution limit of the RGS (e.g. $\\sigma \\lesssim 260$~km~s$^{-1}$ for \\ion{N}{7} He$\\alpha$). In the 5--38~\\AA \\ band, the low and short-on states have power-law continua, while the main-on continuum exhibits excess emission in the $\\sim10$--15~\\AA \\ band. The continuum flux during the short-on is twice as during the low state, a tendency that is followed by most line fluxes as well.  We measure the abundance ratios among C, N, O, and Ne with the following method: we use a power-law emission measure distribution, in conjunction with a grid of XSTAR plasma models, and the $HULLAC$ recombination rates, to fit the narrow line fluxes, and obtain the abundance ratios with the emission measure parameters simultaneously. We perform the measurements for the low and short-on state data separately, and we obtain consistent results with acceptable $\\chi^2$ (Table \\ref{tab:lines}). The enrichment of nitrogen relative to oxygen, of more than four times the solar values, plus the depletion of C and O with respect to Ne, indicate extensive H-burning in a massive star.  The measured abundances require a mechanism for transferring the CNO-processed material onto the HZ~Her envelope. This mechanism is likely linked to the presence of a companion, providing an additional constraint on the evolution of the X-ray binary.  We use spectroscopic analysis and models to set limits on the density and location of the narrow line region. These limits (Fig. \\ref{fig:loclim1}) provide important clues on the nature of the line emission region. We assume thermal and ionization balance to set upper limits on the density, and we use the emission measure derived from the line fluxes to set lower limits on the density. If the line velocity broadening is due to Kepler motion, we may set bounds to the orbital radii. We use the He$\\alpha$ line ratios and UV photoexcitation calculations to set upper limits to the radius enclosing the region.  The low RRC temperatures ($30,000 < T < 60,000$~K) allow us to validate our density upper limits and the photoionization equilibrium models.  The narrow line region may be identified with an accretion disk atmosphere and corona, or with the illuminated face of HZ~Her. The evidence for the disk identification relies on the modeled structure and spectra from a photoionized disk \\cite[]{jimenez}, which agree with the limits set on the density ($10^{13}$--$10^{14}$~cm$^{-3}$). The unresolved velocity broadening indicates the outermost ($r \\sim 5 \\times 10^{10}$ cm) radii of the disk dominate the emission, in agreement with the models, while fluxes of the observed hydrogen and helium-like lines are within a factor of two of the calculations.  Observations from previous missions of the eclipse ingress and egress reveal an emission region which can match the size of the disk and its corona.  The flux variation with orbital phase and with 35~d phase, on the other hand, favors a contribution from the illuminated face of HZ~Her to the narrow line emission.  The variability of the Her X-1 spectrum lends support to the precession of the accretion disk. Regions with contrasting dynamical properties are coming into view at different 35~d phases. Notably, Her X-1 exhibits an ADC-like spectrum during low and short-on states, indicating an edge-on disk, while during main-on states, the spectrum is dominated by the continuum, due to a disk inclination which exposes the pulsar to our line of sight.  We detect an excess of emission centered on the \\ion{O}{7} He$\\alpha$ $i$ line wavelength during the main-on state. If this feature is due to the \\ion{O}{7} He$\\alpha$ $i$ and $r$ lines, their velocity broadening would be in the order of $10^3$~km~s$^{-1}$. Similarly broad, yet double-peaked lines, have to date only been observed from the accreting pulsar in 4U~1626-67 \\cite[]{4u1626}. The observed emission emission feature is noisy, but if it is real, it may be due to emission in the inner accretion disk, near the pulsar magnetopause."
    },
    "0206/astro-ph0206462_arXiv.txt": {
        "abstract": "We have performed TreeSPH simulations of galaxy formation in a standard $\\Lambda$CDM cosmology, including effects of star formation, energetic stellar feedback processes and a meta-galactic UV field, and obtain a mix of disk, lenticular and elliptical galaxies. The disk galaxies are deficient in angular momentum by only about a factor of two compared to observed disk galaxies. The stellar disks have approximately exponential surface density profiles, and those of the bulges range from exponential to $r^{1/4}$ , as observed. The bulge-to-disk ratios of the disk galaxies are consistent with observations and likewise are their integrated $B-V$ colours, which have been calculated using stellar population synthesis techniques. Furthermore, we can match the observed $I$-band Tully-Fisher (TF) relation, provided that the mass-to-light ratio of disk galaxies is $(M/L_I) \\sim$ 0.8. The ellipticals and lenticulars have approximately $r^{1/4}$ stellar surface density profiles, are dominated by non-disklike kinematics and flattened due to non-isotropic stellar velocity distributions, again consistent with observations. ",
        "introduction": "\\vspace{-5mm} The hierarchical Cold Dark Matter (CDM) structure formation scenario has proven remarkably successful on large (cosmological) scales. On galactic scales it has encountered a number of problems, most notably the angular momentum problem, the over-cooling problem, the missing satellites problem and the central cusps problem. Inclusion of the effects of energetic stellar feedback processes in galaxy formation simulations may help to cure a number of these problems as indicated by, e.g., the ``toy'' models of Sommer-Larsen et al. (1999, SLGV). Recently Sommer-Larsen \\& Dolgov (2001, SLD) showed that by going from the CDM structure formation scenario to warm dark matter (WDM) scenarios one can alleviate and possibly even completely overcome the angular momentum problem, and complementary work of Colin et al. (2000) shows this to be the case also for other of the above problems. Fine-tuning of the warm dark matter particle mass to about 1 keV is required, however. In contrast the salient feature about ``conventional'' CDM is that as long as the dark matter particles are much heavier than one keV, the actual particle mass does not matter for structure formation.  We have recently completed a series of considerably more elaborate CDM galaxy formation simulations with very encouraging results: We find that a mix of disk, lenticular and elliptical galaxies can be obtained in fully cosmological ($\\Lambda$CDM), gravity/hydro simulations invoking star-formation, energetic stellar feedback processes and a meta-galactic UV field. These results, together with results on disk gas infall histories and stellar age distributions, hot halo gas properties, global star formation histories etc. are presented in detail in Sommer-Larsen et al. (2002). \\vspace{-5mm} ",
        "conclusions": ""
    },
    "0206/astro-ph0206148_arXiv.txt": {
        "abstract": "We have searched for unresolved X-ray sources in the vicinity of two rich clusters of galaxies: Abell 1995 (A1995) and MS 0451.6-0305 (MS0451), using the \\CHANDRA X-ray observatory. We detected significantly more unresolved sources around A1995 than expected based on the number of X-ray sources to the same flux limit detected in deep \\CHANDRA observations of blank fields. Previous studies have also found excess X-ray sources in the vicinity of several nearby clusters of galaxies using \\ROSAT$ $, and recently in more distant ($z \\approx 0.5$) clusters (RXJ0030 and 3C295) using \\CHANDRA. In contrast, we detect only 14 unresolved X-ray sources near MS0451, which is consistent with the number expected from a cluster-free background. We determine the luminosity functions of the extra sources under the assumption that they are at the distance of their respective clusters. The characteristic luminosity of the extra sources around A1995 must be an order of magnitude fainter than that of the extra sources around RXJ0030 and 3C295. The apparent lack of extra sources around MS0451 is consistent with its greater distance and the same characteristic luminosity as the A1995 sources. Hardness ratios suggest that, on average, the extra sources in A1995 may have harder spectra than those of RXJ0030 and 3C295. These results indicate that different classes of objects may dominate in different clusters, perhaps depending on the formation history and/or dynamical state of the accompanying cluster. ",
        "introduction": "\\label{s:Introduction}    Evidence has accumulated recently that there are more X-ray point sources in the direction of clusters of galaxies than toward cluster-free regions of the sky. Henry \\& Briel (1991), using {\\it ROSAT} PSPC observations, found just about twice as many unresolved sources around Abell 2256 (at $z = 0.06$, Struble \\& Rood 1991) as expected from blank field (no clusters) background observations. The luminosity of these sources, assuming they are at the redshift of the cluster was found to be about $10^{42} \\lum$ or greater (in 0.5 - 2 keV). These sources have high X-ray to optical flux ratios. Some of the sources in A2256 were identified as cluster member galaxies. Henry and Briel also discuss the possibility that the emission from these sources is due to hot gas in galaxies not removed by ram pressure or evaporation, or due to shocks in gas from merging.   Similarly, Lazzati et al. (1998), analyzing ROSAT PSPC images, found an excess number of unresolved X-ray sources in the fields of two nearby clusters: A194 and A1367 ($z =$ 0.018 and 0.022). The spectra of the sources were consistent with thermal bremsstrahlung with $T \\le 2$ keV. Lazzati et al. also found evidence for association of some of these sources with cluster member galaxies, implying luminosities between 0.6 and 6.6 $\\times 10^{41}\\lum$ in the 0.5 - 2 keV band. X-ray emission from hot gas associated with cluster member galaxies had been reported earlier, also based on ROSAT PSPC observations (Grebenev et al. 1995; Bechtold et al. 1983).   Indirect evidence has also been presented for the existence of an excess population of unresolved X-ray sources associated with Abell clusters. Soltan \\& Fabricant (1990) using Imaging Proportional Counter data from the {\\it Einstein Observatory} found excess fluctuations in nearby galaxy clusters which could be explained by assuming the presence of low luminosity sources ($\\approx 4 \\times 10^{41} \\lum$) in clusters with extent less than 1$\\arcmin$. They discuss the possibility that the emission from these sources is due to low luminosity AGNs, or to hot gas in member galaxies. Soltan et al. (1996) found a correlation between the surface brightness of the X-ray background and Abell clusters on scales of a degree, which is much larger than the X-ray emission from the intracluster gas. The characteristic length was found to be about 10 $h^{-1}$ Mpc in radius (where $H_0 = h 100 \\;\\rm Km\\;s^{-1}\\;Mpc^{-1}$), i.e. extra X-ray emission was found around Abell clusters out to about 15 Mpc ($h$ = 0.65). Soltan et al. could not explain the extra X-ray emission based on known sources, or random fluctuations in their number density. They estimated the required number of excess sources to be about 50$\\%$ above the expected number of background sources.   Most recently Cappi et al. (2001), using \\CHANDRA ACIS (Advanced CCD Imaging Spectrometer) observations, found twice as many unresolved X-ray sources in the images of two distant clusters of galaxies, 3C295 ($z = 0.46$ Dressler \\& Gunn 1992), and RX J003033.2+261819 (RXJ0030; $z = 0.5$ Vikhlinin et al. 1998), as expected from a cluster-free background to their flux limits (Giacconi et al. 2001; Mushotzky et al. 2000).  Our main goal in this letter is to present new results on the flux and number distributions of unresolved X-ray sources based on our \\CHANDRA ACIS observations of two rich clusters of galaxies: A1995 and MS0451. We also address briefly the nature of the excess sources.    \\begin{table}[t] \\footnotesize \\begin{center} \\caption{Unresolved X-ray sources detected in A1995.} \\begin{tabular}{lccc} &            &           &        \\\\ \\hline \\hline X-RAY SOURCE  &   COUNTS   &  COUNTS   & R$^a$  \\\\ & 0.5-2 KEV  & 2-10 KEV  & Mag    \\\\ \\hline CXOU J45308.70+580313.8 & 585$\\pm$25 & 179 $\\pm$14 & 18.8 \\\\ CXOU J45305.82+580309.1 & 342$\\pm$19 & 79.1 $\\pm$9.6 & 19.2 \\\\ CXOU J45307.10+580205.8 & 306$\\pm$18 & 74.2 $\\pm$9.2 & 20.8 \\\\ CXOU J45327.65+580339.5 & 166$\\pm$14 & 39.7 $\\pm$7.6 & 20.9 \\\\ CXOU J45305.45+580033.9$^b$ & 69.5$\\pm$8.5 & $<$10         & 10.2 \\\\ CXOU J45246.31+580059.7 & 65.4$\\pm$8.2 & 14.5 $\\pm$4.0 & 20.9 \\\\ CXOU J45317.46+580003.0 & 56.3$\\pm$8.2 & 20.9 $\\pm$5.9 & 20.3 \\\\ CXOU J45233.22+580559.2 & 48.8$\\pm$7.1 & 17.6 $\\pm$4.4 & 19.5 \\\\ CXOU J45229.56+580418.2 & 37.5$\\pm$6.2 & 20.4 $\\pm$4.7 & $>$22 \\\\ CXOU J45230.74+580448.5$^c$ & 34.5$\\pm$6.0 & $<$10           & 16.0 \\\\ CXOU J45233.56+580456.3 & 28.4$\\pm$5.4 & 23.0 $\\pm$5.0 & 22.0 \\\\ CXOU J45324.68+580318.4 & 26.7$\\pm$5.7 & $<$10           & $>$22 \\\\ CXOU J45316.75+575928.6 & 26.6$\\pm$6.1 & $<$10           & $>$22 \\\\ CXOU J45248.65+580255.5$^b$ & 24.3$\\pm$5.3 & $<$10           & 13.2 \\\\ CXOU J45301.83+580005.7 & 22.1$\\pm$5.0 & 11.4 $\\pm$4.0 & $>$22 \\\\ CXOU J45255.19+580056.0 & 20.7$\\pm$4.7 & $<$10           & $>$22 \\\\ CXOU J45244.00+580203.3 & 20.6$\\pm$4.7 & 14.2 $\\pm$4.0 & $>$22 \\\\ CXOU J45315.41+580448.5 & 20.5$\\pm$5.6 & $<$10           & 21.0 \\\\ CXOU J45253.64+580020.3 & 20.1$\\pm$4.6 & $<$10           & $>$22 \\\\ CXOU J45319.11+580134.4 & 17.2$\\pm$4.8 & $<$10           & $>$22 \\\\ CXOU J45313.22+580126.9 & 16.5$\\pm$4.6 & $<$10           & $>$22 \\\\ CXOU J45315.55+580117.4 & 16.3$\\pm$4.9 & $<$10           & 20.3 \\\\ CXOU J45242.50+580159.1 & 15.8$\\pm$4.1 & $<$10           & $>$22 \\\\ CXOU J45228.20+575954.5 & 15.0$\\pm$4.0 & 7.2 $\\pm$3.0 & 21.1 \\\\ CXOU J45231.03+580010.9 & 14.5$\\pm$3.9 & 5.1 $\\pm$2.4 & 21.1 \\\\ CXOU J45235.73+580656.1 & 12.8$\\pm$3.7 & $<$10           & $>$22 \\\\ CXOU J45245.50+580519.5 & 12.2$\\pm$3.6 & $<$10           & $>$22 \\\\ CXOU J45234.49+575904.0 & 11.9$\\pm$3.6 & 14.2 $\\pm$4.2 & $>$22 \\\\ CXOU J45251.93+580046.2 & 10.1$\\pm$3.3 & 12.3 $\\pm$3.7 & 20.9 \\\\ CXOU J45322.04+575858.7 & $<$10           & 25.3 $\\pm$6.4 & $>$22 \\\\  % \\hline\\hline \\end{tabular} \\end{center} $^a$ R band magnitude of optical counterparts/limit if not detected \\\\ $^b$ GSC2 object (\\texttt{http://www-gsss.stsci.edu/gsc/gsc2})      \\\\ $^c$ IRAS Galaxy (F14511+5816)                                      \\\\ \\end{table} ",
        "conclusions": "\\label{s:Discussion}   The derived luminosity functions enable us to estimate the contribution of the unresolved sources to the overall X-ray emission around clusters of galaxies. The derived surface brightness of unresolved X-ray sources in the $8\\arcmin \\times 8\\arcmin$ fields near RXJ0030 and 3C295 is comparable to the surface brightness of the X-ray background ($2.6 \\times 10^{-8}\\flux str^{-1}$ in the 0.5-2 keV band). This enhancement is about 100 times larger than the central enhancement from large scale emission (10 $h^{-1}$ Mpc) found by Soltan et al. (1996). Cappi et al.'s results (cf. their Figure 4) show that these unresolved sources do not extend beyond the clusters much more than about 4$\\arcmin$, as opposed to the large scale component of Soltan et al., which extends out to about 30$\\arcmin$ when scaled to the redshifts of RXJ0030 and 3C295 ($z \\approx$ 0.5). Therefore it is likely that this component contributes only to the compact component found by Soltan et al.   There are a number of possibilities for the nature of these extra sources: cosmic variance, star-burst galaxies, a result of gravitational lensing of background objects, or an enhanced number density of AGNs/QSOs. It is unlikely that the extra sources are due to cosmic fluctuations, which is only at the level of about 20$\\%$-30$\\%$, significantly below the measured factor of two (Cappi et al. 2001). It could be possible, however, that the excess of unresolved sources is due to projection effects, with differences arising from whether we are viewing along or perpendicular to a filament of the cosmic web. Star-burst galaxies are also unlikely to be the sources Cappi et al. found since their X-ray luminosities are about 10-100 times too faint. However, recent \\CHANDRA deep surveys did find some exceptional X-ray bright galaxies at similar redshifts. Gravitational lensing can increase the number of unresolved sources, but only if the \\logn slope is steep enough ($\\ge 0.4$, Croom $\\&$ Shanks 1999; Mellier 1999). Two opposite effects are competing in determining the number of observed sources: lensing magnifies the flux, but it also reduces the field of view behind the gravitational lens. Lensing would need a significantly higher slope in \\logn to explain the large number of extra sources in the field. Refregier \\& Loeb (1997) predict an average reduction of the surface density of faint sources at fluxes less than $10^{-15} \\flux$. Cappi et al. conclude that, as far as spectra and luminosities are concerned, the unresolved sources near RXJ0030 and 3C295 could be AGNs/QSOs associated with the respective clusters.   In Figure~\\ref{F:FIG2} we show the hardness ratios (HRs), $(H-S)/(H+S)$, where $S$ and $H$ are X-ray fluxes in the 0.5-2 keV (soft) and in the 2-10 keV (hard) bands, of unresolved sources in the fields of A1995 (squares), MS045 (triangles) as a function of $H+S$ (our results). As a comparison, we also plot the hardness ratios of RXJ0300 (stars) and 3C295 (diamonds) from Cappi et al. (2001). Points with one sided error bars represent sources not detected either in the hard or in the soft band. The average HR of unresolved sources detected in both soft and hard bands near RXJ0300 and 3C295 are $\\approx -0.5$, while the average HR of unresolved sources near A1995 is slightly harder, about $-0.25$. These sources near A1995 also have lower fluxes than sources near the other two clusters, as previously noted (see section 3).   Comparing our Figure~\\ref{F:FIG2} to Figure 3 of Rosati et al. (2002), which shows the HRs of sources of different types as a function of their luminosities, we conclude that unresolved sources in RXJ0300 and 3C295 with $H+S \\approx 0.004$ cts/sec, corresponding to luminosities of about $10^{44} \\lum$ at the distance of the clusters, would be compatible to the HRs of Type I AGNs (as noted by Cappi et al. 2001). While the faint unresolved sources ($L_* = 5 \\times 10^{41}\\;\\lum$) in the field of A1995 are concentrated around $H+S \\approx 0.0006$ cts/sec, HR $\\approx -0.25$, which falls between normal and star-burst galaxies.   Recent results show that the angular correlation function of X-ray selected AGNs is similar to that of nearby galaxies suggesting that AGNs sample the mass density the same way as galaxies sample (Akylas, Georgantopoulos \\& Plionis 2000), in contrast to optically selected AGNs, which are found to be more frequent in field galaxies (5$\\%$) than in galaxies near clusters (1$\\%$, Dressler Thompson and Shectman 1985; Osterbrock 1960). Therefore we would expect more X-ray selected AGNs in clusters. Since the clustering length of X-ray selected AGNs and nearby galaxies is the same within errors ($\\approx 7\\;h^{-1}$ Mpc, Basilakos 2001; Akylas, Georgantopoulos \\& Plionis 2000; Peebles 1993), we would expect the ratio of total number of galaxies to the number of X-ray selected AGNs, $N_{gal}/N_{AGN}$, to be $\\approx \\langle n_{al}\\rangle/\\langle n_{AGN}\\rangle$ (where $\\langle n_{gal}\\rangle$ and $\\langle n_{AGN}\\rangle$ are the average number densities of galaxies and AGNs in the cluster). However, this effect could only account for about 20$\\%$ of the excess, much less than the a factor of two, which has been found by Cappi et al (2001) and this work.   Our results, that the characteristic luminosities of extra sources are about one order of magnitude different in A1995 vs. 3C295 and RXJ0030, and that there seems to be a difference between their HRs argues against cosmic variance and projection effects. It suggests instead that different class of objects might dominate in different clusters perhaps depending on the formation history and/or the dynamical state of the cluster.   Perhaps the unresolved sources in A1995 belong to a class of starburst galaxies, a result of enhanced star formation due to interactions between infalling groups of galaxies and the intra-cluster gas. This enhanced star formation would lead to an excess of blue galaxies around these areas similar to the Butcher-Oemler effect (Butcher \\& Oemler 1978). A search for a correlation between galaxy color changes around X-ray sources compared to other areas in the cluster could be used to check this possibility.   At present, the exact nature of these objects is not known. Due to limited photon statistic, their spectra could not be determined individually, both low temperature ($\\le 2$ keV) thermal bremsstrahlung and power low spectra can be fitted to their stacked spectra. Revealing the physical properties of these objects would help us to improve our understanding of structure formation, specifically the origin and evolution of the intra-cluster gas, and the effect of merging. Identification of these sources would also help to asses the contamination these sources cause in the interpretation of cluster emission as thermal brems-strahlung from intra-cluster gas. This contamination would result an overestimation of the normalization of the X-ray flux from the cluster and would lead to a systematic error in the determination of the Hubble constant using SZ effect and thermal bremsstrahlung (see for example: Molnar, Birkinshaw and Mushotzky 2002). Follow up observations of the individual sources are necessary to solve this mystery."
    },
    "0206/astro-ph0206476_arXiv.txt": {
        "abstract": "We study a white dwarf model with differential rotation and magnetic field, for which the symmetry axis of the toroidal field, the magnetic axis of the poloidal field, and the principal axis \\( I_{3} \\) coincide permanently; the common axis defined this way is called {}``magnetic symmetry axis''. Furthermore, the magnetic symmetry axis inclines at a small angle \\( \\chi  \\) relative to the spin axis of the model; this angle is called {}``obliquity angle'' or {}``turn-over angle''. Such a model is almost axisymmetric and undergoes an early evolutionary phase of secular timescale, characterized by the fact that the moment of inertia along the spin axis, \\( I_{zz}\\simeq I_{33} \\), is greater than the moments of inertia along the (almost) equatorial axes, \\( I_{11}=I_{22} \\), since rotation and poloidal field (both responsible for the oblateness of the model) dominate over the toroidal field (responsible, in turn, for the prolateness of the model). During this early evolutionary phase, the model suffers from secular angular momentum loss due to weak magnetic dipole radiation activated by the poloidal field. Such an angular momentum loss leads gradually to a situation of dynamical asymmetry with \\( I_{11}>I_{33} \\). However, dynamically asymmetric configurations tend to turn over spontaneously, that is, to rotate about axis with moment of inertia greater than \\( I_{33} \\) with angular momentum remaining invariant. So, the fate of a dynamically asymmetric configuration is to become an oblique rotator and, eventually, a perpendicular rotator. During the so-called {}``turn-over phase'', the turn over angle, \\( \\chi  \\), increases spontaneously up to \\( \\sim 90^{\\circ } \\) on a {}``turn-over timescale'', \\( t_{\\mathrm{TOV}} \\), since the rotational kinetic energy of the model decreases from a higher level when \\( \\chi \\simeq 0^{\\circ } \\) (aligned rotator) to a lower level when \\( \\chi \\simeq 90^{\\circ } \\) (perpendicular rotator). At this lower level the model reaches the state of least energy consistent with its prescribed angular momentum and magnetic field. The excess rotational kinetic energy due to differential rotation is totally dissipated due to the action of turbulent viscosity in the convective regions of the model. In the present paper, we study numerically the so-called {}``turn-over scenario'' (i.e., an evolutionary scenario, which takes into account the turn-over phase) for white dwarf models of several masses, angular momenta, and magnetic fields. ",
        "introduction": "In a recent paper \\citep[hereafter Paper I, and references therein]{ger01}, the turn-over scenario has been studied numerically for rotating magnetic white dwarfs \\citep{ger02}. The turn-over scenario, hereafter TOV scenario, deals with the problem of rotational evolution of a star by taking into account the gradual increase of the turn-over angle, i.e., the angle between the magnetic symmetry axis and the spin axis of the star.  In the TOV scenario, we assume that the differentially rotating white dwarf model is initially almost axisymmetric (i.e., its turn-over angle \\( \\chi  \\) is small: \\( \\chi \\leq 2^{\\circ } \\), say) and undergoes an {}``early evolutionary phase'', during which rotation and poloidal field prevail against the toroidal field, yielding oblate configurations with \\( I_{33}>I_{11} \\) (where \\( I_{33} \\) is the moment of inertia along the magnetic symmetry axis, almost coinciding with the spin axis, and \\( I_{11}=I_{22} \\) are the moments of inertia along the other two principal axes).  However, due to a weak magnetic dipole radiation activated by the poloidal field, the toroidal field becomes gradually more effective, and eventually leads the model to the so-called {}``late evolutionary phase'', during which the model suffers from {}``dynamical asymmetry'': \\( I_{11}>I_{33} \\).  Dynamical asymmetry leads the model to the so-called {}``turn-over phase'', during which the turn-over angle, \\( \\chi  \\), increases spontaneously up to \\( \\sim 90^{\\circ } \\) on a turn-over timescale, \\( t_{\\mathrm{TOV}} \\). The terminal model rotates about its \\( I_{1} \\) axis, coinciding with the invariant angular momentum axis, and occupies the state of least energy consistent with its angular momentum and magnetic field. The excess energy due to differential rotation, defined by the angular velocity component \\( \\Omega _{3} \\) along the spontaneously turning over \\( I_{3} \\) axis (permanently coinciding with the magnetic symmetry axis), is dissipated down to zero due to the efficient action of turbulent viscosity in the convective zone of the model. So, the terminal model does not rotate about its magnetic symmetry axis. Furthermore, it seems difficult for the terminal model to sustain differential rotation along its \\( I_{1} \\) axis, mainly due to the destructive action of the poloidal field. In particular, there is a competition between the efforts of the magnetic stresses to remove rotational nonuniformities, and those of the rotational velocities to bury and destroy the magnetic flux. If the magnetic field and the electrical conductivity have appropriate values (see especially \\S~7 of Paper I), then the magnetic field prevails and removes all the nonuniformities of rotation. So, the terminal model rotates rigidly about its \\( I_{1} \\) principal axis with angular velocity \\( \\Omega _{1} \\).  In Paper I, the aim was to compute the so-called {}``optimal values'' for the angular momentum, \\( L_{xx} \\), the average surface poloidal field, \\( B_{s} \\), and the time evolution parameter, \\( \\delta  \\), under which the model starts its turn-over phase (Paper I, \\S~6). In the present paper, on the other hand, our aim is slightly different. In particular, we shall study several possible turn-over scenarios, corresponding to several indicative values \\( L_{xx} \\) and \\( B_{s} \\). We shall describe in detail our computations in the following sections. ",
        "conclusions": "Computed parameters for the aligned and the perpendicular rotators are given in Tables~\\ref{tab:060_Aligned_constant_for_Bs}--\\ref{tab:132_Perpendicular_variable_for_Bs}. In Tables~\\ref{tab:060_Aligned_constant_for_Bs}, \\ref{tab:090_Aligned_constant_for_Bs}, and \\ref{tab:132_Aligned_constant_for_Bs}, we give parameters regarding the aligned rotators. These parameters are almost independent of the average surface poloidal field, \\( B_{s} \\); namely, the central period, \\( P_{xx} \\), the rotational kinetic energy, \\( T_{xx} \\), the moments of inertia \\( I_{11} \\) and \\( I_{33} \\) along the principal axes \\( I_{1} \\) and \\( I_{3} \\), the average surface toroidal field, \\( \\left\\langle H_{ts}\\right\\rangle  \\), the maximum toroidal field, \\( H_{t\\left[ \\mathrm{max}\\right] } \\), the average toroidal field, \\( \\left\\langle H_{t}\\right\\rangle  \\), and the ratio \\( \\omega ^{-1}_{\\mathrm{e}}=\\Omega _{\\mathrm{c}}/\\Omega _{\\mathrm{e}} \\), where \\( \\Omega _{\\mathrm{c}} \\), \\( \\Omega _{\\mathrm{e}} \\) are the central and equatorial angular velocities, respectively. The corresponding parameters for the perpendicular rotators are given in Tables~\\ref{tab:060_Perpendicular_constant_for_Bs}, \\ref{tab:090_Perpendicular_constant_for_Bs}, and \\ref{tab:132_Perpendicular_constant_for_Bs}. Here, \\( P_{\\mathrm{RR}} \\) is the central period and \\( T_{\\mathrm{RR}} \\) is the rotational kinetic energy of the perpendicular rotator.  Parameters regarding the aligned rotators and varying with the average surface poloidal field, \\( B_{s} \\), are given in Tables~\\ref{tab:060_Aligned_variable_for_Bs}, \\ref{tab:090_Aligned_variable_for_Bs}, and \\ref{tab:132_Aligned_variable_for_Bs}; namely, the corresponding poloidal magnetic parameter, \\( \\beta ^{p}_{*} \\), the surface Alfv\\'en speed, \\( V_{As} \\), the surface Alfv\\'en time, \\( t_{As} \\), the maximum poloidal field, \\( H_{p\\left[ \\mathrm{max}\\right] } \\), and the average poloidal field, \\( \\left\\langle H_{p}\\right\\rangle  \\). The corresponding parameters for the perpendicular rotators are given in Tables~\\ref{tab:060_Perpendicular_variable_for_Bs}, \\ref{tab:090_Perpendicular_variable_for_Bs}, and \\ref{tab:132_Perpendicular_variable_for_Bs}.  General results concerning the turn-over phase are given in Tables~\\ref{tab:060_Turn_Over}, \\ref{tab:090_Turn_Over}, and \\ref{tab:132_Turn_Over}. Parameters tabulated in these tables are the magnetic flux, \\( f \\), the turn-over timescale, \\( t_{\\mathrm{TOV}} \\), the spin-down time rate due to turn-over, \\( \\overset {\\cdot }{P}_{\\mathrm{TOV}} \\), the turn-over timescale in units of the starting central period, \\( N_{xx}=t_{\\mathrm{TOV}}/P_{xx} \\), the present turn-over time, \\( t_{\\mathrm{now}} \\), and the power loss due to turn-over, \\( \\overset {\\cdot }{T} \\). These tables also contain three parameters which are independent of the average surface poloidal field; namely, the rigid rotation amplification ratio, \\( A_{r} \\), the current turn-over angle, \\( \\chi _{\\mathrm{now}} \\), and the current differential rotation strength, \\( F_{r\\left[ \\mathrm{now}\\right] } \\).  Our numerical results reveal that the turn-over timescale, \\( t_{\\mathrm{TOV}} \\), varies from \\( \\sim 0.4 \\) to \\( \\sim 7300 \\) million years, dependent on the mass of the model, the angular momentum, and the average surface poloidal field of the starting model. The turn-over angles \\( \\chi _{\\mathrm{now}} \\), corresponding to the present turn-over times \\( t_{\\mathrm{now}} \\), vary from \\( \\sim 3^{\\circ } \\) to \\( \\sim 90^{\\circ } \\).  An issue of particular interest is the estimated values for the spin-down time rate due to turn-over, \\( \\overset {\\cdot }{P}_{\\mathrm{TOV}} \\), which vary from \\( 4.3\\times 10^{-17}\\, \\mathrm{s}\\, \\mathrm{s}^{-1} \\) to \\( 1.1\\times 10^{-12}\\, \\mathrm{s}\\, \\mathrm{s}^{-1} \\). From columns 1 and 4 of Tables~\\ref{tab:060_Turn_Over}, \\ref{tab:090_Turn_Over}, and \\ref{tab:132_Turn_Over}, it is apparent that, for low values of the surface poloidal field, the turn-over spin-down is negligible, since \\( \\overset {\\cdot }{P}_{\\mathrm{TOV}}\\ll |\\overset {\\cdot }{P}_{\\mathrm{REF}}| \\). However, high values of the surface poloidal field lead to values of \\( \\overset {\\cdot }{P}_{\\mathrm{TOV}} \\) which are comparable to or even greater than \\( |\\overset {\\cdot }{P}_{\\mathrm{REF}}| \\) (Figs.~\\ref{fig:060_P_timerate}, \\ref{fig:090_P_timerate}, and \\ref{fig:132_P_timerate}). For such high values of \\( B_{s} \\), \\( \\overset {\\cdot }{P}_{\\mathrm{TOV}}/|\\overset {\\cdot }{P}_{\\mathrm{REF}}| \\) depends only weakly on the particular values of \\( L_{xx} \\) (roughly speaking, it is independent of \\( L_{xx} \\)). Consequently, in a white dwarf which is now in its turn-over phase, the turn-over effects are negligible if the surface poloidal field is weak; on the other hand, if the surface poloidal field is strong or even moderate, the turn-over effects to the time rate of the central period cannot be neglected."
    },
    "0206/astro-ph0206195_arXiv.txt": {
        "abstract": "We observed 1E 1207.4--5209, a neutron star in the center of the supernova remnant PKS 1209--51/52, with the ACIS detector aboard the {\\sl Chandra} X-ray observatory and detected two absorption features in the source spectrum. The features are centered near 0.7 keV and 1.4 keV, their equivalent widths are about 0.1 keV. We discuss various possible interpretations of the absorption features and exclude some of them. A likely interpretation is that the features are associated with atomic transitions of once-ionized helium in the neutron star atmosphere with a strong magnetic field. The first clear detection of absorption features in the spectrum of an isolated neutron star provides an opportunity to measure the mass-to-radius ratio and constrain the equation of state of the superdense matter. ",
        "introduction": "Although neutron stars (NSs) have been studied extensively  for more than three decades, the properties of the superdense matter in their interiors still remain an enigma.  We know neither the density nor the composition of a NS core. We are not even sure that NSs are indeed composed of neutrons --- for instance, their cores could be pion or kaon condensates or a quark-gluon plasma. A key observational property that would help  understand the true nature of these objects is the mass-radius relation --- if we knew the masses and radii for a sample of NSs, we could  compare them with the predictions of models based on various equations of state of the superdense matter, which  are quite different for NSs of different composition (see Lattimer \\& Prakash 2001 for details). A useful constraint can be obtained via measuring the gravitational redshift $z$ of lines in the spectrum of thermal radiation emitted from the NS surface (atmosphere), which directly gives the mass-to-radius ratio:  $M/R = (c^2/2G) [1-(1+z)^{-2}]$. However, many attempts to detect spectral lines in thermal radiation from isolated (non-accreting) NSs have been unsuccessful. For example, no spectral lines have been found in the spectra of the Vela pulsar (Pavlov et al.\\ 2001), anomalous X-ray pulsar 4U~0142+61 (Juett et al.\\ 2002), and nearby radio-quiet NS RX~J1865--3754 (Burwitz et al.\\ 2001; Drake et al.\\ 2002), despite sensitive observations with the {\\sl Chandra} grating spectrometers. In this Letter we report the first firm detection of absorption features in the spectrum of an isolated neutron star, \\ns.  \\ns, a radio-quiet central source of the SNR PKS~1209--51/52 (also known as G296.5+10.0), was discovered by Helfand \\& Becker (1984) with the {\\sl Einstein} observatory. Mereghetti, Bignami \\& Caraveo (1996) and Vasisht et al.~(1997)  interpreted the \\ros\\ and \\asca\\ spectra of \\ns\\ as blackbody (BB) emission of $T\\simeq 3$~MK from an area with radius $R\\simeq 1.5\\, (d/2\\, {\\rm kpc})$~km. Zavlin, Pavlov \\& Tr\\\"umper (1998) interpreted the observed spectra as emitted from a light-element (hydrogen or helium) atmosphere.  For a NS of mass $1.4~M_\\odot$ and radius 10~km, they obtained a NS surface temperature $T_{\\rm eff}=(1.4$--$1.9)$~MK and a distance $d=1.6$--3.3 kpc, consistent with the distance to the SNR, $d=2.1^{+1.8}_{-0.8}$~kpc from the neutral hydrogen absorption measurements (Giacani et al.\\ 2000).  Zavlin et al.\\ (2000) observed \\ns\\ with the \\chan\\ X-ray Observatory and discovered a period of about 424 ms, which proved that the source is a NS. Second \\chan\\ observation provided an estimate of the period derivative,  $\\dot{P} \\sim (0.7$--$3) \\times 10^{-14}$ s s$^{-1}$ (Pavlov et al.\\ 2002a).  This estimate implies that the characteristic age of the NS, $\\tau_c \\sim 200$--1600 kyr, is much larger than the 3--20 kyr age of the SNR (Roger et al.\\ 1998), while the conventional magnetic field, $B\\equiv 3.2\\times 10^{19} (P \\dot{P})^{1/2}\\, {\\rm G} = (2$--$4) \\times 10^{12}$  G, is typical for a radio pulsar.  Spectral analysis of these \\chan\\ observations, which resulted in the discovery of two absorption features in the NS spectrum, is presented below. ",
        "conclusions": "The discovery of the absorption features in the spectrum of \\ns\\/ provides the first opportunity to measure the mass-to-radius ratio and the magnetic field of an isolated NS.  Measuring $M/R$ is particularly important because it can constrain the equation of state of the superdense matter in the NS interiors, infer the internal composition of NSs, and test the theories of nuclear interactions.  Our analysis of the low-resolution spectra has shown that interpreting the observed features as cyclotron lines requires artificial assumptions, and such features cannot be explained as formed in a hydrogen atmosphere. The energies of the features suggest that they could be associated with the atomic transitions of once-ionized helium in an atmosphere with a strong magnetic field. This interpretation yields a gravitational redshift of about 0.17. To firmly identify the features and measure the gravitational redshift with better accuracy, deep observations with high spectral resolution are needed, such that would be able to resolve potentially multiple lines blended together due to the low spectral resolution of the CCD detector."
    },
    "0206/hep-th0206088_arXiv.txt": {
        "abstract": "{Recently Hollands and Wald argued that inflation does not solve any of the major cosmological problems. We explain why we disagree with their arguments. They also proposed a new speculative mechanism of generation of density perturbations. We show that in their scenario the inhomogeneities responsible for the large scale structure observed today were generated at an epoch when the energy density of the hot universe was $10^{95}$ times greater than the Planck density. The only way to avoid this problem is to assume that there was a stage of inflation in the early universe.}     \\begin{document} ",
        "introduction": "During the last 20 years inflationary theory \\cite{Guth,New,Chaot} has evolved from a problematic hypothesis to an almost universally accepted cosmological paradigm \\cite{book}. It solves many fundamental cosmological problems and makes several predictions that agree very well with the observational data \\cite{Bond}. Despite this fact (or maybe because of it) it has become popular to propose various alternatives to inflation.  In this paper we will consider one such alternative suggested recently by Hollands and Wald \\cite{alt}. The authors admit that their model does not solve or even address the homogeneity, isotropy, flatness, horizon and entropy problems, but they claim that inflation does not do so either. We will examine their claim and explain why we disagree with it. We will use this discussion as an opportunity to emphasize some properties of inflation that may not be widely known.  What Hollands and Wald's model does attempt to explain is the generation of density perturbations with a flat spectrum. However, one cannot justify this mechanism using the standard methods of quantum field theory. Moreover, in this scenario density perturbations on the scale of the present horizon were generated at a time when the energy density of the hot universe was $10^{95}$ times greater than the Planck density. Since nobody knows how to make any calculations at such densities, one must be hard pressed to consider this an alternative to the inflationary mechanism of generation of density perturbations. We will explain that the origin of this problem is directly related to the absence of inflation in the model proposed in \\cite{alt}. ",
        "conclusions": "In this paper we analysed the argument of Hollands and Wald suggesting that inflation does not solve any of the major cosmological problems. Their argument was based on the observation that if our universe were to collapse back to the singularity, it would not deflate. Therefore they argued that it could not inflate on its way to its present state. We do not think that this argument is valid. The inflationary regime is an attractor for solutions for the scalar field during  expansion, but it is a repulsor for the solutions during contraction of the universe. Moreover, the dynamics of inflation are completely irreversible due to particle production after inflation and creation of inhomogeneities during inflation. Therefore the investigation of the time-reversed behaviour of a typical post-inflationary universe tells us almost nothing about the initial conditions that produced the universe. Meanwhile, an investigation performed here and in \\cite{Chaot,book,Creation,Linde:1994wt,Eternal,LLM,Khalat} suggests that initial conditions for inflation in the simplest versions of chaotic inflation are quite natural, and inflation does indeed solve the major cosmological problems.  We also showed that the new mechanism of generation of density perturbations proposed by Hollands and Wald is very problematic. In particular, in their scenario the inhomogeneities responsible for the large scale structure observed today were generated at an epoch when the energy density of the hot universe was $10^{95}$ times greater than the Planck density. This makes all predictions concerning such density perturbations completely unreliable. We have shown that the only way to avoid this problem is to assume that there was a stage of inflation in the early universe.   The authors are grateful to S.~ Hollands, R.~Wald and D. Page for illuminating discussions and to G. Felder for his assistance. The work by L.K. was supported by NSERC and CIAR. The work by A.L. was supported by NSF grant PHY-9870115, and by the Templeton Foundation grant No. 938-COS273.  L.K. and A.L. were also supported by NATO Linkage Grant 97538. V. M. is grateful to Princeton University for the hospitality."
    },
    "0206/gr-qc0206027_arXiv.txt": {
        "abstract": "\\noindent{}It is shown that optical geometry of the Reissner-Nordstr\\\"om exterior metric can be embedded in a hyperbolic space all the way down to its outer horizon. The adopted embedding procedure removes a breakdown of flat-space embeddings which occurs outside the horizon, at and below the Buchdahl-Bondi limit ($R/M=9/4$ in the Schwarzschild case). In particular, the horizon can be captured in the optical geometry embedding diagram. Moreover, by using the compact Poincar\\'e ball representation of the hyperbolic space, the embedding diagram can cover the whole extent of radius from spatial infinity down to the horizon. Attention is drawn to advantages of such embeddings in an appropriately curved space: this approach gives compact embeddings and it distinguishes clearly the case of an extremal black hole from a non-extremal one in terms of topology of the embedded horizon. ",
        "introduction": "This paper introduces a new, simple and convenient way to discuss physics of spacetime near the black hole horizon in terms of the optical geometry embedded in a Poincar\\'e ball. We do not report a progress in developing new physical ideas. Instead, the goal of the present paper is more modest and concentrated on a useful though quite inconventional embedding procedure.  Embedding of curved spaces and spacetimes in a Euclidean space with a higher number of dimensions is a well-known technique, so often used in general relativity that there is no need to recall its usefulness here \\cite{mtw:gravitation}. The optical geometry is also useful, but perhaps less known. In a static spacetime, it is defined by a particular conformal rescaling of the three-dimensional geometry of space (i.e.\\ a hypersurface orthogonal to timelike Killing trajectories) in such a way that light trajectories {\\it in space\\/} are geodesic lines \\cite{abramowicz_etal:1988}. Therefore, the geometry of space can be directly established in terms of measurements based solely on light tracing. This is a rather useful property of the optical geometry that has already been employed by numerous authors for remarkable simplifications of various arguments and calculations, ranging from gravitational wave modes trapped in super-compact stars \\cite{sonego:2000} to origin of the Hawking radiation \\cite{abramowicz:1997}.  Not only are geodesic lines optically straight, but they are also inertially, dynamically and electrically straight and their direction agrees at each point with dynamical, inertial and electrical experiments (e.g., a gyroscope that moves along such a straight line does not precess). These properties follow from the rather remarkable fact that all the relevant equations -- geodesic, Fermi-Walker, Maxwell, Abraham-Lorentz-Dirac, Klein-Gordon -- when written in the 3+1 form of the optical geometry are found to be identical with the corresponding equations in Minkowski spacetime with a scalar field $\\Phi$ (the gravitational potential). For these reasons the geometry of optical space offers simple explanations of several physical effects in strongly curved spacetimes that otherwise could appear unclear or even confusing. See ref.~\\cite{abramowicz:1993} for a general exposition of the properties of optical geometry, and ref.~\\cite{sonego:1998} for a thorough discussion and derivations.  The optical geometry of a given spacetime cannot, in general, be embedded in a Euclidean space all the way to the horizon because of two separate reasons: (i)~the horizon is, obviously, at an infinite distance in the optical geometry based on light tracing; (ii)~optical geometry has a negative curvature near the horizon. In this paper we show, by embedding the optical geometry in a Poincar\\'e ball, how one can avoid both of these difficulties. The Poincar\\'e ball embeddings show clearly the topology of the horizon of the black hole. This feature should be helpful in discussing the role of the horizon topology in the context of the Hawking radiation. Our embeddings illustrate in a striking way the passage from non-extremal to the extremal Reissner-Nordstr\\\"om hole. In particular, the special nature of the extremal Reissner-Nordstr\\\"om hole is clearly manifested in terms of a change in the topology of the horizon. This feature is of interest in the context of supersymmetric theories where the extremal black holes show up as supersymmetric configurations \\cite{horowitz}. ",
        "conclusions": "We have argued in this paper that the Poincar\\'e ball embedding is a natural and useful tool to study global topological properties of spaces with negative curvature, for example the optical space of a black hole geometry.  One particular subject that we plan to study using the Poincar\\'e ball embeddings is the problem of topology of the event horizon in the optical space corresponding to the Reissner-Nordstr\\\"om solution.  In the optical space, the event horizon is always located at infinity, corresponding to the surface of the Poincar\\'e ball in the embedding.  It is interesting to note that while for non-extremal ($Q<M$) black holes the region occupied by the horizon is finite (see fig.~\\ref{fig20}, left panel), in the case of an extremal black hole it shrinks to a single point (fig.~\\ref{fig20}, right panel).   This is because the optical geometry of the extremal black hole is asymptotically flat at both ends.  \\ack VK acknowledges hospitality of SISSA (Trieste), and support from GACR\\,205/00/1685 and 202/02/0735."
    },
    "0206/astro-ph0206060_arXiv.txt": {
        "abstract": "We present CCD $uvby\\beta$ photometry for stars in the nuclei of the young double cluster $h$ and $\\chi$ Persei. We find that the reddening is highly variable through the $h$ Per nucleus, increasing from west to east, with values ranging from $E(b-y) = 0.328\\pm0.022$ in the western part to $E(b-y) = 0.465\\pm0.024$ in the south-east. Towards $\\chi$ Per the reddening is fairly constant, with $E(b-y) = 0.398\\pm0.025$. Both clusters share a common distance modulus of 11.7$\\pm$0.1 mag., and an age of $\\log t = 7.10\\pm0.05$ years. ",
        "introduction": "The precise determination of galactic open clusters main physical parameters plays a central role in the study of the stellar structure and evolution. With accurate photometric data, and once the external variables such as reddening are corrected for, the cluster distances, ages and chemical abundances can be inferred from the study of the photometric colour-magnitude and colour-colour diagrams.  The usual way to obtain the cluster age is by means of isochrone fitting to the main sequence in a colour-magnitude diagram. In the case of young clusters -younger than 50 Myr- isochrone fitting is made difficult by the usual presence of differential reddening across the cluster face, which widens the observed main sequence. Moreover, the presence of emission line stars, like Be o PMS stars, which occupy anomalous positions in the photometric diagrams, additionally contributes to a further main sequence widening. Hence, the fit of a particular isochrone can be a very uncertain process, and it is not difficult to find recent age determinations with very diverging values for a given young cluster.  In a recent paper, Fabregat \\& Torrej\\'on (\\cite{fabregat00}) propose the use of isochrone fitting in the $V_0 - c_0$ plane of the $uvby$ photometric system as an adequate tool to obtain accurate cluster ages. The range of variation of the $c_0$ index along the B-type sequence amounts more than 1 mag., being significantly larger than most common used photometric colours. Moreover, the $c_0$ index is less affected by reddening, and allows an efficient segregation of emission-line stars.  In order to produce accurate and homogeneous dating for a sample of young galactic clusters, we have undertaken an observational programme to obtain CCD $uvby\\beta$ photometry. In this paper we present the first results, related to the clusters $h$ and $\\chi$ Persei.  The double cluster \\object{$h$} and \\object{$\\chi$ Persei} is one of the richest young open clusters in the Galaxy, and also one of the brightest and closest to us. On a clear night, far from the light pollution, it can be easily seen with a naked eye, distinctly shining in the Milky Way between Perseus and Cassiopeia constellations. $h$ and $\\chi$ Persei (\\object{NGC 869} and \\object{NGC 884} respectively) form the nucleus of the broader Per OB1 association.  The published work on $h$ and $\\chi$ Persei is very extensive. A summary of the key papers and discussion on the past work is given by Waelkens et al. (\\cite{waelkens}). Within the literature there is not convergence on the fundamental parameters of the clusters, such as their distances and ages. The discrepancies in the cluster distance moduli are well in excess of 0.5 mag., while there is not agreement about both clusters being of the same or different ages.  Two modern studies based on CCD photometry, presented by Keller et al. (\\cite{keller}) and  Marco \\& Bernabeu (\\cite{marco}), converge in a common distance modulus of about 11.7 mag. for both clusters, but still present controversial results regarding their ages. The former finds a common age of log $t$ = 7.1 years for both clusters and the surrounding population, while the latter claims the existence of at least three different episodes of star formation. To contribute to ascertain this issue is one of the objectives of the present paper. ",
        "conclusions": "We have presented CCD $uvby\\beta$ photometry for stars in the nuclei of the young open clusters $h$ and $\\chi$ Per. We have shown that our photometry is free of systematic effects and well tied to the standard $uvby$ system.  We have obtained the cluster astrophysical parameters from the analysis of the B type star range. The reddening is highly variable through the $h$ Per nucleus, increasing from west to east. Its value ranges from $E(b-y) = 0.328\\pm0.022$ in the western part to $E(b-y) = 0.465\\pm0.024$ in the south-east. Towards $\\chi$ Per the reddening is fairly constant, with $E(b-y) = 0.398\\pm0.025$. Both clusters share a common distance modulus of 11.7$\\pm$0.1 mag., and an age of $\\log t = 7.10\\pm0.05$ years."
    },
    "0206/astro-ph0206256_arXiv.txt": {
        "abstract": "We constrain flat cosmological models with a joint likelihood analysis of a new compilation of data from the cosmic microwave background (CMB) and from the 2dF Galaxy Redshift Survey (2dFGRS). Fitting the CMB alone yields a known degeneracy between the Hubble constant $h$ and the matter density $\\Omega_m$, which arises mainly from preserving the location of the peaks in the angular power spectrum. This `horizon-angle degeneracy' is considered in some detail and shown to follow a simple relation $\\Omega_m h^{3.4} = {\\rm constant}$. Adding the 2dFGRS power spectrum constrains $\\Omega_m h$ and breaks the degeneracy.  If tensor anisotropies are assumed to be negligible, we obtain values for the Hubble constant $h=0.665\\pm0.047$, the matter density $\\Omega_m=0.313\\pm0.055$, and the physical CDM and baryon densities $\\Omega_c h^2 = 0.115 \\pm 0.009, \\Omega_b h^2 = 0.022\\pm 0.002$ (standard rms errors). Including a possible tensor component causes very little change to these figures; we set an upper limit to the tensor-to-scalar ratio of $r<0.7$ at 95\\% confidence.  We then show how these data can be used to constrain the equation of state of the vacuum, and find $w<-0.52$ at 95\\% confidence. The preferred cosmological model is thus very well specified, and we discuss the precision with which future CMB data can be predicted, given the model assumptions. The 2dFGRS power-spectrum data and covariance matrix, and the CMB data compilation used here, are available from {\\tt http://www.roe.ac.uk/{\\tt\\char'176}wjp/}. \\hfill\\break \\hfill\\break {\\bf keywords}: large-scale structure of Universe, cosmic microwave background, cosmological parameters ",
        "introduction": "The 2dF Galaxy Redshift Survey (2dFGRS; see e.g. Colless et~al. 2001) has mapped the local Universe in detail.  If the galaxy distribution has Gaussian statistics and the bias factor is independent of scale, then the galaxy power spectrum should contain all of the available information about the seed perturbations of cosmological structure: it is statistically complete in the linear regime. The power spectrum of the data as of early 2001 was presented in Percival et~al. (2001), and was shown to be consistent with recent cosmic microwave background (CMB) and nucleosynthesis results.  In Efstathiou et~al. (2002) we combined the 2dFGRS power spectrum with recent CMB datasets in order to constrain the cosmological model (see also subsequent work by Lewis \\& Bridle 2002).  Considering a wide range of possible assumptions, we were able to show that the universe must be nearly flat, requiring a non-zero cosmological constant $\\Lambda$.  The flatness constraint was quite precise ($|1-\\Omega_{\\rm tot}| < 0.05$ at 95\\% confidence); since inflation models usually predict near-exact flatness ($\\vert 1 - \\Omega_{\\rm tot} \\vert < 0.001$; e.g. Section 8.3 of Kolb \\& Turner 1990), there is strong empirical and theoretical motivation for considering only the class of exactly flat cosmological models.  The question of which flat universes match the data is thus an important one to be able to answer.  Removing spatial curvature as a degree of freedom also has the practical advantage that the space of cosmological models can be explored in much greater detail.  Therefore, throughout this work we assume a universe with baryons, CDM and vacuum energy summing to $\\Omega_{\\rm tot} = 1$ (cf.  Peebles 1984; Efstathiou, Sutherland \\& Maddox 1990).  In this work we also assume that the initial fluctuations were adiabatic, Gaussian and well described by power law spectra. We consider models with and without a tensor component, which is allowed to have slope and amplitude independent of the scalar component. Recent Sudbury Neutrino Observatory (SNO) measurements (Ahmad et~al. 2002) are most naturally interpreted in terms of three neutrinos of cosmologically negligible mass ($\\simlt 0.05$~eV, as opposed to current cosmological limits of order 2~eV -- see Elgaroy et~al. 2002). We therefore assume zero neutrino mass in this analysis.  In most cases, we assume the vacuum energy to be a `pure' cosmological constant with equation of state $w \\equiv p / \\rho c^2 = -1$, except in Section~\\ref{sec:quin} where we explore $w > -1$.  In Section~\\ref{sec:CMBdata} we use a compilation of recent CMB observations (including data from VSA (Scott et~al. 2002) and CBI (Pearson et~al. 2002) experiments) to determine the maximum-likelihood amplitude of the CMB angular power spectrum on a convenient grid, taking into account calibration and beam uncertainties where appropriate. This compression of the data is designed to speed the analysis presented here, but it should be of interest to the community in general.  In Section~\\ref{sec:models} we fit to both the CMB data alone, and CMB + 2dFGRS.  Fits to CMB data alone reveal two well-known primary degeneracies. For models including a possible tensor component, there is the tensor degeneracy (Efstathiou 2002) between increasing tensors, blue tilt, increased baryon density and lower CDM density.  For both scalar-only and with-tensor models, there is a degeneracy related to the geometrical degeneracy present when non-flat models are considered, arising from models with similar observed CMB peak locations (cf. Efstathiou \\& Bond 1999).  In Section~\\ref{sec:horizon} we discuss this degeneracy further and explain how it may be easily understood via the horizon angle, and described by the simple relation $\\Omega_m h^{3.4} = {\\rm constant}$.  Section~\\ref{sec:quin} considers a possible extension of our standard cosmological model allowing the equation of state parameter $w$ of the vacuum energy component to vary. By combining the CMB data, the 2dFGRS data, and an external constraint on the Hubble constant $h$, we are able to constrain $w$. Finally, in Section~\\ref{sec:cls}, we discuss the range of CMB angular power spectral values allowed by the present CMB and 2dFGRS data within the standard class of flat models. ",
        "conclusions": "Following recent releases of CMB angular power spectrum measurements from VSA and CBI, we have produced a new compilation of data that estimates the true power spectrum at a number of nodes, assuming that the power spectrum behaves smoothly between the nodes. The best-fit values are not convolved with a window function, although they are not independent. The data and Hessian matrix are available from {\\tt http://www.roe.ac.uk/{\\tt\\char'176}wjp/CMB/}. We have used these data to constrain a uniform grid of $\\sim2\\times10^8$ flat cosmological models in 7 parameters jointly with 2dFGRS large scale structure data. By fully marginalizing over the remaining parameters we have obtained constraints on each, for the cases of CMB data alone, and CMB+2dFGRS data. The primary results of this paper are the resulting parameter constraints, particularly the tight constraints on $h$ and the matter density $\\Omega_m$: combining the 2dFGRS power spectrum data of Percival et~al. (2001) with the CMB data compilation of Section~\\ref{sec:CMBdata}, we find $h=0.665\\pm0.047$ and $\\Omega_m=0.313\\pm0.055$ (standard rms errors), for scalar-only models, or $h=0.700\\pm0.053$ and $\\Omega_m=0.275\\pm0.050$, allowing a possible tensor component.  We have also discussed in detail how these parameter constraints arise.  Constraining $\\Omega_{\\rm tot}=1$ does not fully break the geometrical degeneracy present when considering models with varying $\\Omega_{\\rm tot}$, and models with CMB power spectra that peak at the same angular position remain difficult to distinguish using CMB data alone. A simple derivation of this degeneracy was presented, and models with constant peak locations were shown to closely follow lines of constant $\\Omega_m h^{3.4}$. We can note a number of interesting phenomenological points from this analysis:  \\begin{enumerate} \\item The narrow CMB $\\Omega_m-h$ likelihood ridge in Fig.~\\ref{fig:ommvsh} derives primarily from the peak {\\em locations}, therefore it is insensitive to many of the parameters affecting peak {\\em heights}, e.g. tensors, $n_s$, $\\tau$, calibration uncertainties etc. Of course it is strongly dependent on the flatness assumption. \\item This simple picture is broken in detail as the current CMB data obviously place additional constraints on the peak heights. This changes the degeneracy slightly, leading to a likelihood ridge near constant $\\Omega_m h^{3}$. \\item The high power of $h^{3}$ means that adding an external $h$ constraint is not very powerful in constraining $\\Omega_m$, but an external $\\Omega_m$ constraint gives strong constraints on $h$.  A 10\\% measurement of $\\Omega_m$ (which may be achievable e.g. from evolution of cluster abundances) would give a 4\\% measurement of $h$. \\item When combined with the 2dF power spectrum shape (which mainly constrains $\\Omega_m h$), the CMB+2dFGRS data gives a constraint on $\\Omega_m h^2 = 0.1322\\pm0.0093$ (including tensors) or $\\Omega_m h^2 = 0.1361\\pm0.0096$ (scalars only), which is considerably tighter from the CMB alone. Subtracting the baryons gives $\\Omega_c h^2 = 0.1096\\pm0.0092$ (including tensors) or $\\Omega_c h^2 = 0.1151\\pm0.0091$ (scalars only), accurate results that may be valuable in constraining the parameter space of particle dark matter models and thus predicting rates for direct-detection experiments. \\item We can understand the solid contours in Figure~\\ref{fig:ommvsh} simply as follows: the CMB constraint can be approximated as a 1-dimensional stripe $\\Omega_m h^{3.0} = 0.0904\\pm0.0092$ (including tensors) or $\\Omega_m h^{3.0} = 0.0876\\pm0.0085$ (scalars only), and the 2dF constraint as another stripe $\\Omega_m h = 0.20 \\pm 0.03$. Multiplying two Gaussians with the above parameters gives a result that looks quite similar to the fully-marginalized contours. In fact, modelling the CMB constraint simply using the location of the peaks to give $\\Omega_m h^{3.4} = 0.081\\pm0.012$ (including tensors) or $\\Omega_m h^{3.4} = 0.073\\pm0.010$ (scalars only) also produces a similar result, demonstrating that the primary constraint of the CMB data in the $(\\Omega_m, h)$ plane is on the apparent horizon angle. \\end{enumerate}  In principle, accurate non-CMB measurements of both $\\Omega_m$ and $h$ can give a robust prediction of the peak locations assuming flatness. If the observed peak locations are significantly different, this would give evidence for either non-zero curvature, quintessence with $w \\neq -1$ or some more exotic failure of the model.  Using the CMB data to constrain the horizon angle, and 2dFGRS data to constrain $\\Omega_mh$, there remains a degeneracy between $w$ and $h$. This can be broken by an additional constraint on $h$; using $h=0.72\\pm0.08$ from the HST key project (Freedman et~al. 2001), we find $w<-0.52$ at 95\\% confidence. This result is comparable to that found by Efstathiou (1999) who combined the supernovae sample of Perlmutter et~al. (1999) with CMB data to find $w<-0.6$.  In Section~\\ref{sec:cls} we considered the constraints that combining the CMB and 2dFGRS data place on the CMB angular power spectrum. This was compared with the predicted errors from the MAP satellite in order to determine where MAP will improve on the present data and provide the strongest constraints on the cosmological model. It will be fascinating to see whether MAP rejects these predictions, thus requiring a more complex cosmological model than the simplest flat CDM-dominated universe.  Finally, we announce the public release of the 2dFGRS power spectrum data and associated covariance matrix determined by Percival et~al. (2001). We also provide code for the numerical calculation of the convolved power spectrum and a window matrix for the fast calculation of the convolved power spectrum at the data values. The data are available from either {\\tt http://www.roe.ac.uk/{\\tt\\char'176}wjp/} or from {\\tt http://www.mso.anu.edu.au/2dFGRS}; as we have demonstrated, they are a critical resource for constraining cosmological models."
    },
    "0206/astro-ph0206126_arXiv.txt": {
        "abstract": "An overview is presented of the main properties of dark matter haloes, as we know them from observations, essentially from rotation curves around spiral and dwarf galaxies. Detailed rotation curves are now known for more than a thousand galaxies, revealing that they are not so flat in the outer parts, but rising for late-types, and falling for early-types.  A well established result now is that most bright galaxies are not dominated by dark matter inside their optical disks. Only for dwarfs and LSB (Low Surface Brightness galaxies) dark matter plays a dominant role in the visible regions. The 3D-shape of haloes are investigated through several methods, that will be discussed: polar rings, flaring of HI planes, X-ray isophotes. It is not yet possible with rotation curves to know how far haloes extend, but tentatives have been made. It will be shown that the dark matter appears to be coupled to the gas in spirals and dwarfs, suggesting that dark baryons could play the major role in rotation curves. Theories proposing to replace the non-baryonic dark matter by a different dynamical or gravity law, such as MOND, have to take into account the dark baryons, especially since their spatial distribution is likely to be quite different from the visible matter. ",
        "introduction": "\\label{rotation}  At galactic scales, the best tools to probe the dark matter content of the universe are combined HI and H$\\alpha$ or CO rotation curves of spiral galaxies (see the review by Sofue \\& Rubin 2001). The optical rotation curves provide high spatial resolution in the visible disk, and in particular in the center, to trace central mass concentrations, while only the HI gas extend far enough in radius to trace the outer parts, where dark matter is dominating. A lot of progress has been made recently in our knowledge of dark matter content of galaxies, because of large samples observed in 2D Fabry-Perot H$\\alpha$ spectroscopy, and also I-band or near-infrared photometry (Mathewson et al 1992, Schommer et al 1993, Eskridge et al 2000). B-band images of galaxies suffer from extinction, in particular in the center of galaxies, leading to underestimating the stellar disk contribution to the mass, and magnifying the contribution of an hypothetic dark component. Also the mass-to-light ratios are varying more strongly with stellar populations in the blue. This is illustrated by the larger scatter of the Tully-Fischer relation in the blue (e.g. Verheijen 2001).  \\begin{figure} \\resizebox{12cm}{!}{ {\\includegraphics{combesf_f1a.ps}} {\\includegraphics{combesf_f1b.ps}}  } \\caption{ Examples of H$\\alpha$ rotation curves (dots) and their fits with I-band images (full lines); the corresponding M/L ratios is indicated above each panel, from Buchhorn (1993).} \\label{buchhorn} \\end{figure}   \\subsection{Rotation curves of normal spirals}  The new feature resulting from these recent surveys is that for most spiral galaxies, the dark matter is not dominant within the optical disk. Indeed, the 500 rotation curves observed by Mathewson et al. (1992) have been reproduced remarquably well by Buchhorn with mass-to-light ratios constant with radius (e.g. Freeman 1993, and fig \\ref{buchhorn}) and with values compatible with what is known from stellar populations. The fact that the baryonic matter is actually dominant is reflected by the very good fit of all oscillations or \"wiggles\" in the observed rotation curves, corresponding to spiral arms in the disk. A non-baryonic component would not follow the spiral instabilities in the disk, and would have diluted these oscillations in the rotation curves.  This point is related to the maximum disk hypothesis: the latter tries to fit rotation curves in attributing the maximum mass to the disk, compatible to the central part of the curve. Then, keeping the M/L ratio constant with radius, the rotation curve happens to be reproduced quite well over the optical disk, without dark matter. Of course, it is still possible to reduce M/L of the stellar component, and also fit the rotation curve with the addition of a dark matter component. But the peculiar streaming motions features are then less well reproduced (Sackett 1997, Palunas \\& Williams 2000).  Also, the fact that most galaxy disks possess bars, and these bars are rotating rapidly (their cororation is located through resonances in the middle of the disk), pleades in favor of a disk dominated by the visible matter, with a negligible contribution of spherical dark matter;  dynamical friction against the dark matter component would slow down the bars in a few dynamical times (Debattista \\& Sellwood 1998).  \\begin{figure} \\resizebox{6.8cm}{!}{\\includegraphics{combesf_f2a.ps}} \\resizebox{5.2cm}{!}{\\includegraphics{combesf_f2b.ps}} \\caption{ {\\bf (a)} The universal rotation curve of spiral disks at different luminosities (M$_I$); radii are normalised to R$_{200}$, the mean radius containing a mean halo overdensity of 200.  {\\bf (b)} The slope of the rotation curve in the region (0.6-1) R$_{opt}$ versus the rotation velocity V$_{opt}$ at R$_{opt}$. From Persic et al. 1996). } \\label{persic96} \\end{figure}  The bottom line is therefore that dark matter is only needed at large radii, in the HI-21cm extensions. The fact that the rotation curve is flat in the outer parts, while it is no longer attributable to the stellar disk or bulge, has been called the conspiracy. Why does the velocity due to the spherical non-baryonic dark matter coincide exactly to that of the stellar component?  In fact, rotation curves are not all flat, depending on their morphological types (Casertano \\& van Gorkom 1991): early-type galaxies have rotation curves that begin to fall down, while late-type and dwarfs have not yet reached their maximum velocity at the last observed radius.  When compiling 1100 rotation curves, and normalising them to their exponential radial scales, Persic et al. (1996) found that spiral discs (once the contribution of the bulge is removed in galaxies), may have an universal rotation curve, only determined from their total luminosity (cf fig \\ref{persic96}). At high luminosities, there is no or only a slight discrepancy between the observed rotation curve and that contributed by the luminous matter, while a larger dark matter component is required at low luminosities: the dark-to-luminous mass ratio scales inversely with luminosity (fig \\ref{persic96}). The halo core radius is comparable to the optical radius, but shrinks for low luminosities.     However, to draw these conclusions, Persic et al. (1996)  assumed that there is a constant ratio between the end-radius of the visible disk (R$_{23.5}$), and the exponential characterictic radius, or in other words, that all disks have the same shape. This is not quite true, as emphasized by Palunas \\& Williams (2000). the latter authors have carried out a detailed study of 74 galaxies, where 2D Fabry-Perot H$\\alpha$ spectroscopy exist (Schommer et al 1993) and I-band photometry. Very good fits of the rotation curves are obtained without dark matter, out to R$_{23.5}$, with a constant M/L. They conclude that mass traces light, in particular since the surface brightness profiles of the various galaxies present pronounced differences. The small number of galaxies with a poor fit have strong non-axisymmetric structures (bars and strong spiral arms). The resulting I-band M/L = 2.4 $\\pm$ 0.9 h$_{75}$, is compatible with normal stellar populations. This indicates that the dark matter is not dominant within optical disks, or is perfectly coupled to the visible matter.   Already this fact is contradictory to expectations from CDM scenarios. CDM halo profiles are centrally concentrated, and numerical simulations predict that the dark matter dominates inside spiral disks. For example in a galaxy of the mass of the Milky Way, $\\Lambda$CDM simulations predict three times more dark matter than is observed (Steinmetz \\& Navarro 2000). On the contrary, this fact is in agreement with MOND hypothesis.  \\subsection{Rotation curves of dwarfs and LSB}  The relative importance of dark matter is increasing towards late types and dwarf irregular galaxies are completely dominated by dark matter. They are ideal tools to probe theories of dark matter, since the uncertainties on the stellar mass-to-light ratio has negligible influence on the derived radial matter profile. For the prototype of these dwarfs, DD0154, the rotation curve is well determined until 15 optical scale lengths; the HI gas component is more massive than the stellar disk (Carignan \\& Beaulieu 1989).  The derived radial profile of dark matter in dwarfs is not peaked towards the center, since the rotation curves are slowly rising. This is one of the main problems for the $\\Lambda$CDM theories: the radial distribution is predicted by simulations to be highly peaked, with a cusp, or density following a power-law of slope -$\\alpha$ = -1.5 (Navarro, Frenk \\& White 1997, Fukushige \\& Makino 97). Observed rotation curves points towards no cusp, but cores (Moore 1994, Dalcanton \\& Bernstein 2000). According to Burkert \\& Silk (1997), this problem could ony be solved by the introduction of baryonic dark matter inside the optical disk, with a mass several times the visible mass, and with a similar radial distribution.  However, there are still uncertainties in the mass-to-light ratios, and the rotation curves are not fully sampled in all dwarf galaxies available, so that it might be still difficult to conclude for all of them (Swaters 1999, van den Bosch \\& Swaters 2000). New models of dark matter have been proposed to solve precisely this problem, self-interacting dark matter with a non-zero cross-section (Spergel \\& Steinhardt 2000), but many new problems then appear. Other mechanisms have been proposed, such as stellar feedback, to reduce central densities of CDM (Navarro et al 1996, Binney et al 2001); but this mechanism has very low efficiency, as soon as the galaxy is more massive than 10$^7$ M$_\\odot$ (Mc Low \\& Ferrara 1999).  Low Surface Brightness galaxies (LSB) are also dominated by dark matter; they can be dwarfs, but also massive galaxies, with a large amount of HI gas. Their rotation curves are also good constraints for dark matter models. Again, they are incompatible with the cuspy profiles predicted for $\\Lambda$CDM, but can be fitted with models where matter follows light, although with too large mass-to-light ratios (de Blok et al 2001). ",
        "conclusions": "Dark haloes at galactic scales are now constrained by more precise rotation curves. Bright spiral galaxies are not dominated by dark matter in their optical disks. The dark-matter/visible mass ratio is a function of luminosity and surface brightness. Dwarf and LSB galaxies are the best laboratories for dark matter studies since they are dominated by unseen matter down to their central regions. The derived radial profile of dark matter is not centrally concentrated and presents no cusp as predicted in the CDM scenario.  The 3D shape of haloes is still badly constrained. Polar rings are often self-gravitating and there are some clues that their potential is flattened along the polar plane. The HI plane flaring method depends strongly on the assumed truncation radius of the dark matter component.  Observations have shown however that haloes are oblate and galaxy potential axisymmetric in their planes.  Statistical galaxy-galaxy lensing might bring some progress in the determination of shape and radial extension of dark matter haloes.  The formation of tidal tails in galaxy interactions is a good test of the shape of their potential. Simulations have shown that only galaxies dominated by their visible matter, or with their halo truncated outside their optical disk, were able to form tails corresponding to observations. Such simulations should be explored within the MOND hypothesis.  Most baryons are dark, according to primordial nucleosynthesis and CMB anisotropies. These baryons could be present in the form of cold molecular clouds in the outer parts of galaxies, with a H$_2$/HI surface density ratio of about 10, as suggested by rotation curves. This reservoir of gas in the outer parts account for galaxy evolution, that requires fresh replenishment of gas for star formation, and explains the evolution of morphological types along the Hubble sequence: late-types have a much larger proportion of dark matter than early-types, while secular evolution (through bars and spirals), and interactions/mergers tend to progressively transform late-type galaxies in early-type ones.  In galaxy clusters, the baryonic matter is almost all visible in the form of hot X-ray gas. The distribution of the dark with respect to visible matter, which increases with radius at galactic scales, and then decreases with radius at cluster scale, might raise strong constraints in all modified gravity/dynamics theories. For all these scenarii, the spatial distribution of baryonic dark matter is a fundamental element to consider."
    },
    "0206/astro-ph0206310_arXiv.txt": {
        "abstract": "s{ A brief overview on the theory and observations of relativistic particle populations in clusters of galaxies is given. The following topics are addressed: (i) the diffuse relativistic electron population within the intra-cluster medium (ICM) as seen in the cluster wide radio halos and possibly also seen in the high energy X-ray and extreme ultraviolet excess emissions of some clusters, (ii) the observed confined relativistic electrons within fresh and old radio plasma and their connection to cluster radio relics at cluster merger shock waves, (iii) the relativistic proton population within the ICM, and its observable consequences (if it exists), and (iv) the confined relativistic proton population (if it exists) within radio plasma. The importance of upcoming, sensitive gamma-ray telescopes for this research area is highlighted.} ",
        "introduction": "Even though the study of relativistic particle population is more than three decades old, it has recently received a significant increase in attention by various researchers. Here, a brief and, therefore, incomplete and personally biased overview of this field is provided. A guide through the lines of argumentation is given by Fig. \\ref{fig:diag}, which sketches the main dependencies of the components of the theory and their observational consequences. This figure is explained in the following.  The main energy sources of the relativistic particle population in clusters are outflows from galaxies (galactic winds, radio jets) and/or the energy released in accretion on galaxy clusters. The first sources can directly eject relativistic particles into the ICM \\cite{1977ApJ...212....1J,1993ApJ...406..399G}, whereas the latter produce shock waves and turbulence, which can accelerate particles via the Fermi mechanisms \\cite{1987A&A...182...21S,1997A&A...321...55D,1997MNRAS.286..257K,1998AA...332..395E,2001MNRAS.320..365B,2001ApJ...559...59M,2001ApJ...562..233M}. Also the termination shocks of galaxy winds were proposed as shock acceleration sites \\cite{1996SSRv...75..279V}. A different source of relativistic particles may be the annihilation of certain dark matter particles \\cite{2001ApJ...562...24C}.  The relativistic particles loose energy via various radiative and non-radiative processes, allowing to measure or constrain their spectral energy distribution observationally. The most important loss channels are discussed in the following.  \\begin{figure} \\begin{center} \\psfig{figure=diagram.new.eps,width=0.85\\textwidth} \\end{center} \\caption{\\label{fig:diag} Basic theory building blocks covering most of the proposed scenarios for the production and maintenance of relativistic particle populations in galaxy clusters and their observational consequences.  For details see text.} \\end{figure} ",
        "conclusions": ""
    },
    "0206/astro-ph0206382_arXiv.txt": {
        "abstract": "The origin of rovibrational \\H2 emission in the central galaxies of cooling flow clusters is poorly understood. Here we address this issue using data from our near-infrared spectroscopic survey of 32 of the most line-luminous such systems, presented in the companion paper by Edge et al.~(2002).  We consider excitation by X-rays from the surrounding intracluster medium (ICM), UV radiation from young stars, and shocks. The v=1-0 K-band lines with upper levels within $10^{4}$\\K~of the ground state appear to be mostly thermalised (implying gas densities $\\approxgt 10^{5}$\\pcm), with the excitation temperature typically exceeding 2000\\K, as found earlier by Jaffe, Bremer \\& van der Werf~(2001). Together with the lack of strong v=2-0 lines in the H-band, this rules out UV radiative fluorescence.  Using the CLOUDY photoionisation code, we deduce that the \\H2 lines can originate in a population of dense clouds, exposed to the same hot ($T \\sim 50000$\\K) stellar continuum as the lower density gas which produces the bulk of the forbidden optical line emission in the H$\\alpha$-luminous systems. This dense gas may be in the form of self-gravitating clouds deposited directly by the cooling flow, or may instead be produced in the high-pressure zones behind strong shocks. Furthermore, the shocked gas is likely to be gravitationally unstable, so collisions between the larger clouds may lead to the formation of globular clusters. ",
        "introduction": "In the companion paper by Edge et al.~(2002) (hereafter paper 1) we presented H and K-band spectra of 32 of the most H$\\alpha$ line-luminous cooling flow (CF) central cluster galaxies (CCGs) in the {\\em ROSAT} Brightest Cluster Sample (Crawford et al.~1999). One or more of the rovibrational \\H2 v=1-0 S(0)--S(7) lines were detected in 23 systems; one or both of [FeII]$\\lambda\\lambda$1.258,1.644 were seen in 14 systems, and there is also evidence for higher excitation \\H2 lines (e.g. v=2-1 S(1), S(3)) and some coronal lines ([Si VI], [Si XI], [S XI] and [Ca VIII]). This sample quadruples the number of CF CCGs with \\H2 detections, and builds on work by Jaffe \\& Bremer~(1997), Falcke et al.~(1998), Donahue et al.~(2000) and Jaffe, Bremer \\& van der Werf~(2001). Now that warm ($\\sim 1000-2500$\\K) molecular hydrogen is known to accompany ionized material in the core of a flow, our aim in this paper is to shed light on the associated excitation mechanisms.  Jaffe \\& Bremer~(1997) found that the \\H2 emission in CF CCGs is too luminous to be simply due to material passing through $\\sim 2000$\\K -- the temperature at which it is collisionally excited -- whilst cooling from $\\sim 10^{7}$\\K~(the well-known `H rec' problem). Based on a high v=1-0 S(1)/H$\\alpha$ ratio, they implicated collisional excitation by suprathermal electrons liberated by X-ray photoionization deep within the cold, molecular core of the line-emitting clouds; the X-rays were hypothesized to originate in the surrounding intracluster medium (ICM). A similar model was advanced by Wilman et al.~(2000) for Cygnus A, which exhibits \\H2, H recombination and [FeII] emission out to radii of 5\\kpc. Flux ratios involving the v=1-0 S(1), S(3) and S(5) lines deviate markedly from the predictions of LTE, suggesting that non-thermal excitation is important. If the latter is provided by hard X-rays from the obscured quasar nucleus, the implied \\H2 mass within 5\\kpc~of the nucleus is $\\sim 10^{10}$\\Msun, accounting for $10^{9}$\\yr~of mass deposition from the CF.  From their HST NICMOS narrow-band imaging of the K-band \\H2 emission in three CF CCGs, Donahue et al. deduced that the emission in NGC 1275 is principally due to the AGN (as did Krabbe at al.~2000), but that the more spatially-extended emission in A2597 and PKS 0745-191 is probably due to UV radiation from young stars. They based this conclusion on the close morphological correspondence between the \\H2 and optical line emission, and on a comparison of the \\H2/H$\\alpha$ ratios with predictions for shock models, X-ray heating by the ICM, and UV fluorescence by stars and AGN.  Most recently, Jaffe, Bremer \\& van der Werf~(2001) presented UKIRT CGS4 K-band spectra of 7 well-known cooling flows. They observed that the relative strengths of the v=1-0 lines agree well with the predictions of LTE, but that there are deviations amongst the higher vibrational states. The thermal nature of the line production implies that the molecular gas is overpressurised by 2 to 3 orders of magnitude with respect to the X-ray and optical line emitting components. They sought to explain this using an ablative or `rocket' model of the gas system in which X- or UV-radiation heats the surface layers of a cold, gravitationally bound molecular cloud to 2000\\K, at which temperature the material ceases to be gravitationally bound and expands into space, ionizes, and emits the optical lines. They deduced ionized to molecular line ratios which are lower than those of starburst regions, indicating that alternative heating mechanisms are necessary.  Of relevance to the analysis presented in this paper are the results of the CO survey for cold molecular gas in CF CCGs presented by Edge (2001), many of whose targets are also in our UKIRT survey. He detected CO emission in 16 of these CFs, consistent with $10^{9-11.5}$\\Msun~of molecular gas at $\\sim 40$\\K~(assuming a standard CO:\\H2 conversion factor). The implied mass of cool \\H2 correlates better with the H$\\alpha$ luminosity than with the global X-ray mass deposition rate of the CF, suggesting that young stars are warming a population of molecular clouds. Since the physical conditions in the material probed by the K-band \\H2 lines are likely to occupy an intermediate regime between the optical and CO line-emitting components, they offer the potential for further investigation into the link between the CO and H$\\alpha$ emission. Indeed, we showed in paper 1 that there is a good correlation between the \\H2~mass inferred from the CO emission and the 1-0~S(1) line luminosity, with the exceptions of the strong radio sources Cygnus A and PKS~0745-191. ",
        "conclusions": "Our UKIRT spectra demonstrate that the central galaxies of cooling flow clusters with H$\\alpha$ luminosities above $10^{41}$\\ergps~exhibit rovibrational \\H2 line emission at the level of 0.01--0.1 H$\\alpha$. The relative strengths of the \\H2 lines imply that the emission is thermally excited in dense gas ($n \\approxgt 10^{5}$\\pcm) at temperatures $\\sim 2000$\\K, and is thus overpressurised by 2--3 orders of magnitude with respect to the optical emission line gas and the X-ray ICM. The emission could originate in a population of dense clouds heated by the young stellar populations which are known to exist in such systems and which can reproduce the optical forbidden line spectrum when incident on lower density gas. These dense clouds may be self-gravitating or confined in the high-pressure regions behind strong shocks. The fact that the high and low density clouds ultimately have a similar origin and are excited by the same radiation field accounts for the similar morphologies of the \\H2 and optical emission line images found by Donahue et al.  Future observations at higher spectral resolution could address the kinematics of the \\H2 emission for comparison with those of the CO emission (most of whose line widths are known to within 50\\kmps~and lie in the range 100--300\\kmps; Edge~2001), which may originate at a greater depth within the same dense clouds. Deeper spectra of an enlarged sample of low H$\\alpha$ luminosity systems ($<10^{41}$\\ergps) should also be obtained to establish at what level they emit \\H2. Fig.~4 tentatively suggests that any \\H2 in these systems may be produced by a different mechanism. Our models also open the possibility of detecting some of the purely rotational \\H2 lines in the mid-infrared with forthcoming instruments such as Michelle on the UKIRT and Gemini telescopes. Observations of the 0-0~S(3)/1-0~S(1) ratio, for example, would provide a useful constraint on the variation of temperature with depth over the HI--\\H2 transition region within the clouds."
    },
    "0206/astro-ph0206511_arXiv.txt": {
        "abstract": "In 3-d SPH simulations of the coalescence of a quark star with a pseudo-Newtonian black hole all of the quark matter is quickly accreted by the black hole. The Madsen-Caldwell-Friedman argument against the existence of quark stars may need to be re-examined. ",
        "introduction": "An astrophysical argument has been invoked against the existence of quark stars in our Galaxy: in the coalescence of a quark star with a comparably compact object a huge number of small fragments of quark matter is expected to be ejected from the binary and to pollute the galactic environment\\cite{madsen88,cf91}, precluding formation of glitching neutron stars \\cite{alpar87}.  We have set out to test this expectation against actual simulations of the coalescence process.  In a first attempt to model quark stars in smooth particle hydrodynamics (SPH), we have already carried out strictly Newtonian simulations of the black hole coalescence of stars modeled with an equation of state (e.o.s.) appropriate to self-bound quark matter \\cite[henceforth paper I]{lkn01} and compared the outcome against previously published results \\cite{lk99a,lk99b,kl98,l00,l01} of analogous simulations of the coalescence of a black hole and a polytrope, taken to represent a neutron star.  We have found significant differences between the two sets of simulations. Although the star was disrupted, to a degree, and a disk of matter formed around the black hole in each of the two cases, for the quark star system we have found no clear evidence of mass ejection in those Newtonian simulations, to the limit of our resolution.  Here, we report the results of 3-d hydrodynamic simulations of the coalescence of a quark star moving in a pseudo-potential (e.g., Paczy\\'nski and Wiita 1980) modeling salient features of general relativistic motion around black holes.  The coalescence is over much more quickly than in the Newtonian case. The black hole swallows the quark star in one gulp. ",
        "conclusions": ""
    },
    "0206/astro-ph0206277_arXiv.txt": {
        "abstract": "Using {\\sl RXTE} and {\\sl ASCA} data, we investigate the roles played by occultation and absorption in the X-ray spin pulse profile of the Intermediate Polar PQ Gem. From the X-ray light curves and phase-resolved spectroscopy, we find that the intensity variations are due to a combination of varying degrees of absorption and the accretion regions rotating behind the visible face of the white dwarf. These occultation and absorption effects are consistent with those expected from the accretion structures calculated from optical polarisation data. We can reproduce the changes in absorber covering fraction either from geometrical effects, or by considering that the material in the leading edge of the accretion curtain is more finely fragmented than in other parts of the curtain. We determine a white dwarf mass of $\\sim$1.2 using the {\\sl RXTE} data. ",
        "introduction": "\\label{intro}  Magnetic cataclysmic variables (MCVs) can be split into two groups: those in which the magnetic field of the accreting white dwarf is strong enough ($\\ga$ 10MG) to synchronise its spin period with that of the binary orbital period -- the polars, and those with a magnetic field insufficiently strong to achieve this synchronisation -- the intermediate polars (IPs).  The MCV PQ Gem is unusual in that it exhibits characteristics of both groups: it shows a strong soft X-ray component ($kT\\sim$50eV) (Duck et al. 1994), it is polarised in the optical/IR wave-bands (Potter et al. 1997; Piirola, Hakala \\& Coyne 1993) and has an estimated magnetic field strength of $\\sim$8--21 MG (V\\\"{a}th, Chanmugam \\& Frank 1996; Potter et al. 1997; Piirola et al. 1993) -- all of which are characteristics of polars. On the other hand it shows a spin period of 833.4 sec (Mason 1997) and an orbital period of 5.19 hrs (Hellier, Ramseyer \\& Jablonski 1994) which are typical of IPs. PQ Gem can therefore be thought of as the first true ``intermediate'' polar (Rosen, Mittaz \\& Hakala 1993).  PQ Gem has been observed using several X-ray satellites ({\\sl ROSAT}, {\\it Ginga}, {\\sl ASCA} \\& {\\sl RXTE}). The X-ray light curves show a prominent modulation on the spin period, in particular, a pronounced dip in the light curve which is thought to be due to an accretion stream obscuring the main emission region on the surface of the white dwarf.  Mason (1997) made a study of the then available X-ray data to determine an accurate ephemeris for PQ Gem based on timings of the dip. However, a detailed study of the spectral information contained in the {\\sl ASCA} data was not undertaken.  This paper is primarily targeted at reaching a greater understanding of the interplay between the emission sites and absorption which produces the observed modulation of the X-ray light curves. This is achieved through analysis of the  {\\sl ASCA} spectral  data with supporting evidence from the hard X-ray {\\sl RXTE} light curves.  The mass of the white dwarf is calculated from the {\\sl RXTE} spectral data using the stratified accretion column model of Cropper et al. (1999, subsequently CWRK). A fuller report is available in James (2001). ",
        "conclusions": "In section~\\ref{phasespect} we find that the spin phase modulation can be modelled well by a variation in the partial covering fraction of a neutral absorber at low energies and in the normalisation at higher energies. Before we discuss this further, we briefly examine whether the modulation at both high and low energies can both be explained by a tall accretion shock.  In the case where the shock has significant vertical extent, the lower (cooler) part of the shock could be obscured by the limb of the white dwarf at certain spin phases.  The light curves will then have a larger amplitude at low energies than at high energies (e.g. Allan, Hellier \\& Beardmore 1998). The height of the accretion column, $H$, can be estimated from the relationship \\begin{displaymath} H=5.45\\times 10^8\\dot{M}^{-1}_{16}f_{-2}M_{WD}^{3/2}R_{WD}^{1/2} \\end{displaymath} (Frank, King \\& Raine 1991). Using the lower limit of the fractional area and the corresponding $\\dot{M}$ derived in \\S \\ref{doccult} (cf Table \\ref{tab:mdot}) gives an upper bound of $H$=5.4$\\times$10$^7$cm or 0.14 R$_{WD}$. However, the observational evidence suggests that PQ Gem has a significant magnetic field (\\S \\ref{intro}) which implies that the shock height will be lower than this (e.g. Cropper et al. 1999) due to cyclotron cooling. This suggests that this mechanism is not the cause of the variation. However, the accretion regions are likely to be sufficiently structured (e.g. Potter et al 1997) so that such a scenario cannot be excluded. On the other hand, because the spectral variations can be well modelled by a variation in covering fraction and normalisation, we go on to consider this explanation in more detail.   \\begin{table*} \\begin{center}\\small \\begin{tabular}{p{18mm}p{21mm}p{13mm}p{26mm}p{13mm}p{13mm}p{15mm}} $\\dot{m}$            & WD mass        &$\\chi^2_\\nu$ (dof)&unabsorbed flux          & luminosity        & $\\dot{M}$         & fractional   \\\\ &                &                  &(0.001-100.0keV)         &                   &                   & area          \\\\ (g cm$^{-2}$s$^{-1}$)& $(M_{\\sun})$   &                  &(ergs cm$^{-2}$ s$^{-1}$)& (ergs s$^{-1}$)   & (g s$^{-1}$)      &               \\\\ \\hline 0.5                  & 1.20(1.13-1.26)&0.91(52)          &6.3$\\times 10^{-11}$     &1.2$\\times 10^{32}$&3.0$\\times 10^{15}$&3.1$\\times 10^{-3}$\\\\ 1.0                  & 1.21(1.16-1.24)&0.89(52)          &6.4$\\times 10^{-11}$     &1.2$\\times 10^{33}$&2.9$\\times 10^{15}$&1.6$\\times 10^{-3}$\\\\ 5.0                  & 1.22(1.15-1.29)&0.86(52)          &6.6$\\times 10^{-11}$     &1.3$\\times 10^{33}$&2.8$\\times 10^{15}$&3.4$\\times 10^{-4}$\\\\ \\end{tabular} \\caption{The effect fixing $\\dot{m}$ to a range of values in the CWRK model during the spectral fitting. Details are given in \\S \\ref{doccult}.} \\label{tab:mdot} \\end{center} \\end{table*}   \\subsection{Spin Pulse Modulation at Higher Energies} \\label{dspinmod_h}  It is evident from an inspection of the 8--25 keV spin-phased {\\sl RXTE} light curve (Figure 2) and the spin-phased normalisation in the {\\sl ASCA} spectroscopy (Figure 5) that they are broadly similar. As this emission is expected to be at most weakly beamed, this indicates that the variation in intensity is caused by changes in visibility of the X-ray emitting region as the white dwarf rotates. The phase of maximum emission therefore corresponds to the phase of maximum visibility of the emitting region. There is possibly a slight phase shift between the two curves, with a clear peak at $\\phi$ = 0.2 in the {\\sl RXTE} light curve, and a more extended maximum around 0.2 $\\leq$ $\\phi$ $<$ 0.4 in the {\\sl ASCA} normalisation, but it is unclear given the uncertainties in the normalisations whether this is significant. Spin phases 0.2--0.4 also correspond to the phases at which the accretion region in the upper hemisphere is seen most close to face on (Potter et al 1997, figure 9). This suggests that this is principle cause of the spin pulse modulation at higher energies.  \\subsection{Spin Pulse Modulation at Lower Energies} \\label{dspinmod_l}  At lower energies, the maximum in the covering fraction of the absorber at $\\phi$ = 0.0 is the major effect on the soft X-ray light curve, while the local increase at $\\phi$ = 0.6 combined with the decrease in normalisation at this phase to cause the second minimum (\\S \\ref{phasespect}). Potter et al. (1997) find that accretion occurs preferentially along field lines which thread the disc ahead of the accreting pole (their figure 12). Material accreting along these field lines can therefore be identified as the source of the absorption at phase 0.0. Similarly, at $\\phi$ = 0.6, in the absence of absorption by the disc (cf \\S \\ref{doccult}), the local increase in covering fraction is likely to be caused by absorbing material along field lines which intersect the line of sight to the lower pole (Potter et al. 1997, figure 9).  Figure \\ref{fig:phspct} indicates that the change in absorption can be explained {\\it entirely} by a change in covering fraction, with the column density of the absorber remaining constant. With reference to figure 9 of Potter et al. (1997), at $\\phi$ = 0.0, all of the field lines between the disc and the accreting pole will intersect the line of sight. At $\\phi$ = 0.2, those field lines to the leading part of the accretion region will no longer intersect the line of sight, while by $\\phi$ = 0.4, only field lines to the trailing part of the accretion region will do so. If the column density through accreting field lines is $\\sim 5\\times 10^{22}$ cm$^2$, the observed variation in covering fraction in Figure \\ref{fig:phspct} can be reproduced.  The accretion flow may be more finely fragmented along field lines feeding the leading edge of the accretion region. This would be expected because finely fragmented material is threaded by the magnetic field more easily than the larger denser inhomogeneities (e.g. Wickramasinghe 1988).  From considerations of the packing fraction, the line of sight through more finely fragmented material is less likely to pass between gaps in the flow than in the case of the larger inhomogeneities. This effect would  reproduce the high covering fraction at $\\phi$ = 0.0, and its subsequent decrease towards $\\phi$ = 0.5.  \\subsection{Accretion Model}  The accretion scenario we are proposing fits neither the standard occultation model of King \\& Shaviv (1984) nor the accretion curtain model of Rosen, Mason \\& C\\'{o}rdova (1988). The orientation of the accretion region at spin phase maximum is as predicted by the occultation model, but with absorption effects modifying the light curve significantly. In this it is similar to the ``weak field/fast rotator'' model of Norton et al. (1999) with the symmetry of the accretion curtain about the magnetic axis modified by the leading field lines preferentially stripping material from the inner margin of the disc. The high (among IPs) magnetic field in PQ Gem and relatively high inclination and low dipole offset ($\\sim 60^{\\circ}$ and $\\sim 30^{\\circ}$, Potter et al. 1997) ensure that material on accreting field lines travels far enough out of the orbital plane to pass through the line of sight to the accretion region, causing the observed absorption effects. In this it differs substantially from EX Hya which possesses a magnetic field $<$ 1MG.  \\subsection{Occultation by the Accretion Disc?} \\label{doccult}  Finally, we check whether the accretion disk can extend close enough to the white dwarf for it to have an affect on the X-ray light curves by obscuring the lower emission region.  Using the Ghosh \\& Lamb formulation (Li, Wickramasinghe \\& Rudiger 1996) the radius to the truncated inner edge of the accretion disc, $r_A$, is given by \\begin{eqnarray*} r_A = 0.52\\mu_{WD}^{4/7}(2GM)^{-1/7}\\dot{M}^{-2/7} \\end{eqnarray*} where $\\mu_{WD}$ is the magnetic moment of the white dwarf.  The magnetic moment can be estimated from the relationship $B = \\mu/r^3$. Using the fits to the {\\sl RXTE} data and the model of CWRK we found a best fit to the white dwarf radius, $R_{WD}$, of $3.8\\times 10^8$ cm.  Hence, with a magnetic field strength, $B_{WD}$, of 15MG (Piirola et al. 1993, V\\\"{a}th et al. 1996, Potter et al. 1997), $\\mu_{WD} = 9.6\\times10^{32}$ G cm$^{3}$. The accretion luminosity, \\begin{eqnarray*} L_{acc} = GM\\dot{M}/R_{WD} \\end{eqnarray*} where $M$ is the mass of the white dwarf and L$_{acc}$ is emitted mostly in the X-ray energy band, enables estimation of the accretion rate, $\\dot{M}$.  The unabsorbed spectral model from the analysis of the integrated spectrum (\\S \\ref{intspct}) extrapolated for the energy range 0.001--100.0keV gives the X-ray flux at $3.1\\times10^{-10}$erg cm$^{-2}$ s$^{-1}$ which, taking the distance to PQ Gem of 400pc (Patterson 1994), gives a luminosity of 6.0$\\times10^{33}$erg s$^{-1}$ and hence a mass transfer rate rate, $\\dot{M}$, of $1.4\\times10^{16}$g s$^{-1}$. This is typical for IPs (Warner 1995).  The resulting $r_A = 1.3\\times 10^{10}$ cm or $\\sim34R_{WD}$ may be too large given that an estimate of the distance to the first Lagrangian point is $\\approx$ 200 R$_{WD}$ (Plavec \\& Kratochvil 1964) (the main uncertainty in $r_A$ is in the magnetic moment $\\mu_{WD}$). Nevertheless it does indicate that with a system inclination of $60\\degr$ (Potter et al. 1997), the line of sight to the white dwarf surface is likely to be clear of the accretion disc at all spin phases, and this is unlikely to contribute to the cause of the covering fraction variation.  \\subsection{The size of the Accretion Region}  Using the radius of the white dwarf, the specific mass accretion rate, $\\dot{m}$, (2.3 g cm$^{-2}$s$^{-1}$ cf Table \\ref{tab:massfit}) and $\\dot{M}$ determined above we can derive a fractional accretion area, $f$, of 9.0$\\times 10^{-3}$.  This is within the normal expectation of 0.001$\\la f \\la$ 0.02 for an IP (Rosen 1992). However, from the {\\sl RXTE} data $\\dot{M}$ =2.9 $\\times 10^{15}$ g s$^{-1}$ which implies a fractional area of only 7.1 $\\times 10^{-4}$. Refitting the {\\sl RXTE} data to the model in which $\\dot{m}$ was fixed at 0.5, 1.0 and 5.0 g cm$^2$ did not give a significant adverse effect on the fit ($\\chi^2_\\nu$=0.91, 0.89, 0.86, respectively). Table \\ref{tab:mdot} gives the implied fractional area and $\\dot{M}$ for these $\\dot{m}$ as well as the results from the spectral fitting. For a low specific mass accretion rate we find that the implied fractional area is consistent with the lower limit determined using previous observations.  \\subsection{The Mass of the White Dwarf} \\label{dmass}  Previous determinations of the mass of the white dwarf in PQ Gem using an emission model fitted to {\\it Ginga} data (Cropper et al. 1998, 1999) gave estimates $\\geq$ 1.1 $M_{\\sun}$. In the case of the IP XY Ari, Ramsay et al. (1998) found that there was a good correspondence between the estimates given by this model and those from eclipse mapping.  In our work it is found that the {\\sl RXTE} data gives estimates which are very much better constrained than those made with the {\\sl ASCA} SIS data.  The estimate from our {\\sl RXTE} data of $M_{WD}$ = 1.21 (1.16-1.28) $M_{\\sun}$ corresponds well to that of 1.21 ($>1.08) M_{\\sun}$ obtained with {\\it Ginga} data. Although this appears to be unusually high Ramsay (2000) found that the white dwarf in magnetic CVs were biased towards higher masses compared to isolated white dwarfs."
    },
    "0206/astro-ph0206041_arXiv.txt": {
        "abstract": "We made CO and HCN simultaneous observations of lenticular galaxies, NGC 404, NGC 3593 and NGC 4293, and detected HCN emission in NGC 3593 and NGC 4293 as well as CO in all the galaxies. The $I_{\\rm HCN}$/$I_{\\rm CO}$ ratios were $0.025 \\pm 0.006$ and $0.066 \\pm 0.005$ in NGC 3593 and NGC4293, respectively, which are comparable to the late-type spiral galaxies. The average of the $I_{\\rm HCN}$/$I_{\\rm CO}$ ratios at the center of 12 nearby spiral galaxies including late-type was $0.055 \\pm 0.028$. The line profiles of CO and HCN emission showed different shape in both galaxies. The HCN peaks were not at the systemic velocity of these galaxies, while the CO peaks were near the systemic velocity. These results suggest that the fraction of the dense molecular gas is high around the center in these galaxies. ",
        "introduction": "Since it has been known that even early-type galaxies contain an interstellar medium (ISM), the properties of the ISM have been examined. These investigations have revealed the similarities and dissimilarities of the properties of the ISM in early and late-type spiral galaxies. For example, lenticular galaxies contain almost an order of magnitude less molecular gas than late-type spiral galaxies (Wiklind, Henkel 1989; Thronson et al. 1989), and the molecular gas concentrates in the central region in early-type spirals (Taniguchi et al. 1994; Young et al. 1995). The molecular to atomic gas mass ratio is larger in early-types than late-types (Young, Knezek 1989; Sage 1993; Casoli et al. 1998). On the other hand, it has been shown that star formation efficiency (SFE) derived from star formation rate (SFR) and molecular gas mass are similar in early and late-type spirals (Wiklind, Henkel 1989; Thronson et al. 1989; Rownd, Young 1999).  About the central activity, it has been suggested that there exists a difference in the star formation activity and AGN with Hubble type. Ho et al. (1997b) show that H$\\alpha$ luminosity of nuclei is significantly enhanced in early-type galaxies from their spectroscopic survey. Alonso-Herrero and Knapen (2001) also found the same tendency from the analysis of archival HST/NICMOS H-band and Pa$\\alpha$ data. They suggest that bars affect the star-forming activity, especially in early-type spirals. Ho et al. (1997a) show that AGNs are found predominantly in luminous, early-type galaxies, while H\\,{\\footnotesize II} nuclei prefer less luminous, late-type systems. Furthermore, Lei et al. (2000) indicate from their analysis of a magnitude-limited sample of LINERs that the intensity of AGN activity increases with decreasing star forming contribution from late-type spiral galaxies to early-type ones. They suggest an evolutionary connection between AGNs and starbursts in LINERs.  Dense molecular gas is thought to play an important role for star formation. Moreover, concentration of dense molecular gas has been found in some of the Seyfert nuclei (e.g., Kohno 1998). Therefore, it is very interesting to investigate from observations of tracers of dense molecular gas whether there is the difference of density of molecular gas in the central region between early and late-type galaxies. The number of early-type galaxies in which tracers of dense molecular gas have been detected is still very small (Henkel, Wiklind 1997; Kohno et al. 2001). Since the molecular gas in early-type galaxies is expected to concentrate in the central region, it is preferred to observe nearby galaxies with a small beam size to make the beam-filling factor large. Therefore, we made CO and HCN simultaneous observations of nearby early-type galaxies using the 45-m telescope at the Nobeyama Radio Observatory (NRO). Because of the difference of the critical density for collisional excitation of CO (a few 10$^{2}$ cm$^{-3}$) and HCN ($> 10^{4}$ cm$^{-3}$) lines, the ratio of the intensity of these lines gives a rough estimation of density of the molecular gas. We selected three lenticular galaxies (NGC 404, NGC3593 and NGC4293) in which strong CO emission have been detected (Wiklind, Henkel 1989). ",
        "conclusions": "\\subsection{NGC 404} NGC 404 is an S0 galaxy with central star-forming activity and classified as a LINER (Ho et al. 1997c). We adopted the center position determined by the optical data (Palumbo et al. 1988). The position is 1$''$ south from the compact UV source in the nucleus found by Maoz et al. (1995) and a few arc second east from the center of the CO map of Wiklind and Henkel (1990) obtained with Onsala 20-m telescope. Wiklind and Henkel (1990) found that the CO emission is not centrally peaked. The CO distribution shows an arc-like feature that corresponds to the dust lane seen in the optical photograph, and whose size in the CO map is about 40$''$. Our 19$''$ beam for HCN corresponds to 220 pc and does not cover the whole arc-like structure seen in CO emission, while the bright UV sources observed by Maoz et al. (1995) including the components from young stars near the nucleus are fully covered. Although Wiklind and Henkel (1990) assert that the distance of the galaxy is 10 Mpc, we adopted 2.4 Mpc derived from the systemic velocity from Tully (1988).  Figure 1a shows a CO spectrum measured by our observations. Its line shape is similar to the profile at about 10$''$ east from the center in the map of Wiklind and Henkel (1990) rather than that in the center. This may be due to the pointing errors of their and our measurements. Since our adopted center position is a few arc second east from their center postion and the pointing errors are estimated about 5$''$ and 4$''$ for ours and theirs, respectively, our result does not conflict with theirs. The column density of H$_{2}$ derived using the conversion factor of $2.3\\times 10^{20}$ cm$^{-2}$ [K km s$^{-1}$]$^{-1}$ (Strong et al. 1988) is 47 $M_{\\solar}$ pc$^{-2}$. We did not detect HCN emission from this position. The 2-sigma upper limit of $I_{\\rm HCN}$/$I_{\\rm CO}$ ratio ($I \\equiv \\int T_{\\rm MB} dV$) is 0.042, which is still comparable with late-type spirals (Sorai et al. 2002; Kohno 1998). Since we observed only one position, we can not deny that we missed a clump of dense molecular gas. The CO peak found by Wiklind and Henkel (1990) is shifted by ($+10''$, $-10''$) from our observed position. At least, it is necessary to observe the CO peak, since some galaxies show a HCN peak that corresponds to CO peaks out of the nuclei (e.g., NGC3593, see section 3.2).  \\subsection{NGC 3593} NGC 3593 is an S0/a galaxy that has a counterrotating stellar disk (Bertola et al. 1996). Hunter et al. (1989) showed that this galaxy is under going a burst of star formation and that the star-forming regions are concentrated in the small central region. Corsini et al. (1998) found that the star-forming regions consist of a filamentary pattern with a circumnuclear ring. The high resolution CO maps of the central region of NGC 3593 show three peaks (Sakamoto et al. 2000; Garc\\'ia-Burilo et al. 2000). The central peak is located at the galactic center and other two peaks seem to be associated with the circumnuclear star-forming ring. Our 19$''$ beam for HCN, which corresponds to 510 pc assuming that the distance of NGC 3593 is 5.5 Mpc (Tully 1988), covers most of the H$\\alpha$ ring whose diameter is about 17$''$. We adopted the radio continuum peak (Condon 1987) as the center and detected both CO and HCN emissions (figure 1b). The CO profile and intensity are consistent with that obtained with the IRAM 30-m telescope (Wiklind and Henkel 1992). The nearest point to our observed position is (10$''$, 0$''$) in the IRAM map. The column density of H$_{2}$ is estimated to be 199 $M_{\\solar}$ pc$^{-2}$ using a conversion factor of $2.3\\times 10^{20}$ cm$^{-2}$ [K km s$^{-1}$]$^{-1}$. On the other hand, our HCN spectrum shows a different profile from that detected by Henkel and Wiklind (1997) with the IRAM 30-m telescope. Our spectrum is narrower ($\\Delta V$ = 100 km s$^{-1}$) and stronger ($T_{\\rm peak}$ = 20 mK) than theirs ($\\Delta V > 200$ km s$^{-1}$; $T_{\\rm peak} <$ 10 mK). The difference of the beam size and the observed position may be the cause of the difference of the profiles. Unfortunately, because their observed position was not stated, we can not check the difference between our observed position and theirs.  It is interesting that the CO and HCN spectra show different shape. Since we observed these lines simultaneously, the difference is real and not due to the pointing error. The HCN emission was detected in the velocity range only 500 - 600 km s$^{-1}$, while the CO emission was detected in the velocity range 500 - 800 km s$^{-1}$ at the zero intensity level. The difference is apparent also in the FWHM. The FWHM of the HCN profile is about 50 km s$^{-1}$, while that of the CO profile is about 130 km s$^{-1}$. The peak velocity of the CO profile is 605 km s$^{-1}$ which is near the systemic velocity of the galaxy. On the other hand, the HCN spectrum does not have a peak at the velocity. The peak velocity of the HCN profile is 540 km s$^{-1}$. The $I_{\\rm HCN}$/$I_{\\rm CO}$ ratio is 0.025 $\\pm$ 0.006. The ratio of $T_{\\rm MB}$ in the velocity range where HCN was detected is higher ($\\sim 0.07$).  It is apparent from the velocity field in Garc\\'ia-Burilo et al. (2000) that we detected HCN emission from the western peak only. This result means that this part has a higher fraction of dense molecular gas than in the other region within our 19\" beam. This region coincides with the site where the CO(2 -- 1)/CO(1 -- 0) ratio is high and the most vigorous star-forming region (Wiklind, Henkel 1992; Corsini et al. 1998). It seems to be reasonable that the star-forming activity is high there because of the high fraction of the dense molecular gas.  \\subsection{NGC 4293} NGC 4293 is an S0/a galaxy and classified as LINER (Ho et al. 1997c). There is a dust lane near the center (B\\\"oker et al. 1999). H$\\alpha$ emission concentrates on the nucleus except for another peak about 40$''$ east from the center (Koopmann et al. 2001). Our 19$''$ beam for HCN, which corresponds to 1600 pc at the adopted distance of 17 Mpc, covers the dust lane and the central H$\\alpha$ peak.  We detected CO and HCN emission toward the center of NGC 4293 (figure 1c). The $I_{\\rm HCN}$/$I_{\\rm CO}$ ratio is 0.066 $\\pm$ 0.005 and the ratio of $T_{\\rm MB}$ at the peak of HCN is higher ($\\sim 0.1$), which are highest among three galaxies we observed. The CO profile is quite similar to the CO(2 -- 1) spectra in Wiklind and Henkel (1989). The column density derived from the $I_{\\rm CO}$ is 145 $M_{\\solar}$ pc$^{-2}$ assuming the conversion factor of $2.3\\times 10^{20}$ cm$^{-2}$ [K km s$^{-1}$]$^{-1}$. As seen in NGC 3593, the CO and HCN spectra show a difference in the shape. The HCN emission is weak at the systemic velocity and has a double peak, while the CO spectrum has a single peak near the systemic velocity of this galaxy. Simple interpretation of this result is that dense molecular gas traced by HCN distributes around the center and has a ring-like structure whose size is less than our beam size, while less dense molecular gas traced by CO concentrates in the center. The dense gas is expected to associate with the star-forming activity in the center. Usui et al. (2001) found a correlation between turn-over radius of the rotation curve from rigid rotation to flat rotation and the radius of the active star-forming region in the central regions of early-type spirals. Since the rotation curve of this galaxy rises steeply near the center (Rubin et al. 1999), the diameter of the ring-like structure, if it exists and associates with the star-forming activity in the central region, must be small. The velocity separation of the double peak of HCN is about 100 km s$^{-1}$ in radial velocity. From the rotation curve, the velocity separation corresponds to a radius of a few arc second, or of about 200 pc. Therefore, if the dense molecular gas exists in the turn-over radius, the radius is expected to be about 200 pc. These features resemble the structures seen in NGC 3593, although it seems that the distribution of the dense gas in the ring of NGC 3593 is not symmetric with respect to the center. It is highly desirable to get high-resolution images of CO and HCN to confirm the structures. Although the large inclination angle of this galaxy is a disadvantage, the velocity information will give us some clues about the spatial structure. Furthermore, high-resolution image of tracers of star forming regions are required to examine the relation between dense molecular gas and star formation there.  \\subsection{Comparison with later type galaxies} As compared with the previous survey of HCN including late-type galaxies, the $I_{\\rm HCN}$/$I_{\\rm CO}$ ratio we obtained (table 2) is comparable to the late-type galaxies. Figure 2 shows the correlation between $I_{\\rm CO}$ and $I_{\\rm HCN}$, including the data toward the center of other galaxies measured with the NRO 45-m telescope from the literature for the comparison (Sorai et al. 2002; Kohno 1998; Kohno et al. 2002). The range and average of the distance of the galaxies are 1.8 - 24.1 Mpc and 8 Mpc, respectively, which are comparable with our sample (2.4 - 17 Mpc, 8.3 Mpc). We used the data of the same telescope, since the beam size and the intensity calibration method are the same.  The lenticular galaxies are distributed in lower side of $I_{\\rm CO}$ in figure 2. It means that the column density of molecular gas averaged by our 19$''$ beam is lower in the lenticular galaxies. Since the galaxies plotted here are CO bright galaxies in each morphological type, this result indicates that the upper side of the dispersion of the column density of molecular gas at the central region is higher in later type galaxies. The rate of concentration of molecular gas in the center may be one of the reasons of this trend. Taniguchi et al. (1994) showed that three S0 galaxies out of four have significantly small scale lengths of radial distribution of the molecular gas. The shorter scale length of the column density of molecular gas in early-type galaxies will make the averaged surface density within our 19$''$ beam lower, even if there is no difference in the surface density within much smaller area.  The average of the $I_{\\rm HCN}$/$I_{\\rm CO}$ ratio of all the data plotted in figure 2, except for the upper limits, is 0.055 $\\pm$ 0.028 (1$\\sigma$). On the other hand, the best fit with power-law is $I_{\\rm HCN} \\propto I_{\\rm CO}^{0.9}$. Since the sample number is still small, the result is not conclusive. Especially, when we take into account the upper limit in the lower side, the power seems to be larger than 1. It must be important to check the relation to examine the mechanism of star formation. The star formation scenario based on the large-scale gravitational instability seems to do well to explain the relation between SFR and the surface density of interstellar gas (Kennicutt 1989). Most of the recent results of the observational studies about the relation between SFR and molecular gas show a non-linear relation and a power-law slope of 1.2 -- 1.4 (Kennicutt 1998; Rownd, Young 1999). However, the mechanism that makes stars from the molecular gas in large molecular complexes formed by gravitational instability is still an open question. Because star formation should begin from creation of dense molecular gas that can be traced by HCN, we may be able to examine different phases in the process of star formation, namely, from diffuse molecular gas to dense molecular gas and from dense molecular gas to stars, from the relations between $I_{\\rm CO}$ and $I_{\\rm HCN}$ and between $I_{\\rm HCN}$ and SFR, comparing with the relation between $I_{\\rm CO}$ and SFR.  The different shape of HCN and CO spectra in the early-types in figure 1 is very interesting. For late-type galaxies, the difference is not prominent (e.g., Sorai et al. 2002). Although the beam size of HCN is slightly larger than that of CO, it must be difficult to attribute the difference of the profiles to the difference of the beam size. If the difference attributed to the beam size, the HCN emission must exit just out of the CO beam. Our results imply that in NGC 3593 and NGC 4293 the fraction of dense molecular gas is higher around the center and the dense gas may be associated with a ring-like structure. Pogge and Eskridge (1993) found that a ring of H\\,{\\footnotesize II} regions is fairly common feature in S0 galaxies. The dense molecular gas must be associated with the star-forming activities in these galaxies as seen in NGC 3593. As mentioned in the previous section, the location of the star-forming region seems to be correlated with the turn-over radius of the rotation curve. Actually, the radius of the ring-like structure of CO and H$\\alpha$ in NGC 3593 corresponds to the turn-over radius of the rotation curve (Corsini et al. 1998; Garc\\'ia-Burillo et al. 2000). This relation is explained by viscosity of molecular clouds (Icke 1978; Fukunaga 1983). Namely, in the region of the differential rotation, the molecular clouds lose angular momentum due to shear motions, while shear motions do not occur in the region of the rigid rotation. As a result, the molecular gas accumulates in the turn-over radius. The same relation has been observationally found in late-type galaxies (Fukunaga 1984; Nishiyama et al. 2001). If the molecular gas is accumulated by this mechanism, it is expected that the molecular gas concentrates in a narrower region in early-type galaxies than late-type ones, since rotation curve arises more steeply in early-type galaxies than late-type. Therefore, we speculate that the density of the molecular gas might become extremely high in such a region in early-type galaxies and active star formation might occur there, although, of course, it must depend on the total amount of the accumulated gas. Since the sample galaxies we selected are CO bright galaxies, observations of galaxies in wider range of CO brightness are required to avoid a selection effect. Further observations with an interferometer are also required to know the distribution of molecular gas, especially dense molecular gas, and relations with star-forming activity and dynamics of molecular gas."
    },
    "0206/astro-ph0206331_arXiv.txt": {
        "abstract": "{ In this work we investigate the evolution of the X-ray emission of a cluster of single young massive stars with different metallicities. We have considered the X-ray contribution originated by the diffuse gas heated by the mechanical energy released by stellar winds and SN explosions as well as the X-ray contribution from SN remnants.  The resulting ionizing spectrum (i.e. $\\lambda < 912$ \\AA) has been used to compute the expected intensity of the nebular \\ion{He}{ii} $\\lambda$4686 \\AA. The observational ratio \\ion{He}{ii}/H$\\beta$ could be reproduced by the models assuming that a fraction of the mechanical energy produced by the star-formation episode is reprocessed by interaction with the ISM as soft X-ray radiation, contributing to the He ionization.  However, the discreteness of the stellar populations affects the ionizing flux and may be responsible for the observed dispersion of the ratio.  We have finally used the synthesis models to estimate the contribution of circumnuclear star-forming regions to the multiwavelength energy distribution in Active Galactic Nuclei, finding that the UV to soft X-ray continuum in many Seyfert 2 galaxies seems to be dominated by star-formation processes. ",
        "introduction": "In the last years convincing observational evidences have been collected about the presence of starburst regions in or around active Seyfert~2 nuclei \\citep[ and references therein]{Hecketal97,GDetal98}. It has been found that most of their UV light originates in (circum-)nuclear star-formation sites; the possible connection between the nuclear activity and the properties of these starbursts is still a matter of debate. According to the unified scheme of Seyfert galaxies, the active nucleus should be hidden by an opaque torus in the case of Seyferts~2, explaining why the collected UV light is dominated by young, massive stars. The low degree of contamination by the active source in the UV provides detailed information about the properties of the star-formation processes. Moreover, extrapolating to the radio -- X-ray ranges the emission associated with the starbursts, it should be possible to disentangle the fractional contribution of both sources (the starburst and the active nucleus) at different energy ranges.  Multiwavelength evolutionary synthesis models normalized to the observed UV emission would be the ideal tool to perform this analysis.  With these ideas in mind we started some years ago a program to extend our evolutionary synthesis models (\\citealp{AKS89}, Paper I, \\citealp{MHK91}, Paper II, hereafter MHK, \\citealp{CMH94}, Paper III, hereafter CMH) to high-energy ranges (soft and hard X-rays, $\\gamma$-rays). The predictions of $\\gamma$-ray emission were presented in \\citet{Cervetal00}, and this paper is devoted to the X-ray emission. It became evident from the beginning that to properly reproduce the X-ray emission, both the contribution of evolved binary systems and the heating of the diffuse interstellar gas by the release of mechanical energy (by stellar winds and supernova explosions), had to be considered.  Nevertheless, while High Mass X-ray binaries contribute essentially to the hard X-ray range \\cite[above few keV,][]{vBetal99,VBV2000}, the soft X-ray emission is dominated by very hot diffuse gas, heated by the release of mechanical energy from the starburst \\citep{Hecketal95, Hecketal96, Dicetal96, SS98a, SS98b, SS99, SSS01}.  We will therefore discuss in this paper the predicted soft X-ray emission, and its effects on other observables, like the relative intensity of Hydrogen and Helium emission lines. The contribution of binary systems and evolved sources to the hard X-ray emission will be discussed in a forthcoming paper.  Additionally, the statistical dispersion due to the discreteness of actual stellar populations as presented in \\citet{CLC00,CVGLMH01,Cetal01} is particularly important in the high energy domain, where the number of effective sources is lower. This computed statistical dispersion allows to perform a better comparison with real systems and to evaluate the statistically relevant sources in each energy range.  We present in Sec.~\\ref{sec:model} our evolutionary synthesis model and show how the X-ray emission and its associated statistical dispersion have been computed. In Sec.~\\ref{sec:output} we show the predictions on soft X-ray emission. In Sec.~\\ref{sec:lines} we explore the effects of the soft X-ray contribution on the He~{\\sc ii} nebular emission line.  In Sec.~\\ref{sec:agn} we compare our predictions with observational data from star-forming and Seyfert galaxies, aiming to disentangle the relative contribution of starbursts to the global energy budget of Seyferts. We finally summarize in Sec.~\\ref{sec:summ} our conclusions. ",
        "conclusions": "\\label{sec:summ}  In this work we have explored the X-ray emission originated in a star-forming region with only single stellar populations.  The mechanical energy injected into the interstellar medium by stellar winds and supernova explosions will heat the diffuse gas to very high temperatures, and will be finally reprocessed into soft X-ray emission. We have found that the X-ray emission observed in starburst galaxies can be well explained assuming that a moderate fraction of the mechanical energy is finally reprocessed into X-ray emission. The rest of the mechanical energy released leads to the expansion of the gas, creating bubbles and gas flows at galactic scales. An interesting implication is that the age deduced from kinematical studies will result systematically lower than the one obtained from the analysis of the emission line spectrum if the heating of the gas and the X-ray emission is not taken into account.  Additionally, the inclusion of X-ray transient systems, like SN explosions and Be/X-ray binaries showing bursts of X-ray emission, may lead to some degree of variability in the high energy emission of starburst galaxies.  We have computed the expected intensity of the nebular \\ion{He}{ii} $\\lambda$4686 \\AA, including the additional ionization of the gas by the soft X-ray emission originated in the diffuse gas.  The observational values of the \\ion{He}{ii}/H$\\beta$ ratios can be reproduced by our models assuming moderate efficiencies (about 20\\%) in the reprocessing of mechanical energies into X-ray emission. Alternatively, the mechanical energy reprocessing could not be a continuous process. This energy released by the massive stars could be ``accumulated'' in the form of accelerated gas flows, which would release all this accumulated energy only when the outflowing gas interacts with the static interstellar medium. Under this scenario the soft X-ray emission originated by the shocked gas would provide enough additional ionizing power to explain the observed \\ion{He}{ii}/H$\\beta$ ratios.  We have also shown that such ratios have a high intrinsic dispersion and that a deeper statistical study is necessary to investigate the source of the nebular \\ion{He}{ii} $\\lambda$4686 \\AA ~line.   We have compared the predicted soft X-ray emission with observational values for a sample of star-forming and Seyfert galaxies, aiming to disentangle the contribution of star-formation episodes to the total energy budget of low activity galaxies. We have found that while the high energy emission of Seyfert~1 galaxies is clearly above the predictions of starburst models, both the UV and soft X-ray emissions of many Seyfert~2 galaxies are apparently associated mostly to the (circum-)nuclear starburst episodes known to be present in these objects. The active source in Seyfert~2 galaxies would dominate therefore only at harder X-rays."
    },
    "0206/astro-ph0206107_arXiv.txt": {
        "abstract": "The understanding of pulsar radio emission demands a close look at all regions of the huge parameter space present for pulsar observations. In this review we concentrate on the space given by the range of observed pulse periods, spanning four orders of magnitude. A comparative study of the emission properties of millisecond pulsars and normal pulsars, both at radio and high energy frequencies, promises to shed light on the still poorly understood emission theory. ",
        "introduction": "The search for a theoretical framework that is able to explain the radio emission of pulsars, has been largely unsuccessful. This is due to the difficulty to include all the different phenomena and time scales observed in radio pulsars into a working model. It is essential to isolate the bigger picture, and one way of doing this, is to enlarge the parameter space. Rather than making the picture even more complicated by more observations that are difficult to explain, the aim must be to study the emission properties by pushing the known parameter space to new boundaries to provide solid and obvious constraints.  The parameter space relevant for pulsar radio emission is large. Obvious parameters are the period and its increase, determining important values like magnetic field, or potential drop above the surface. Further parameters are time resolution, frequency coverage in the radio, but also from the radio regime to high energies, eventually simultaneously, and sensitivity. A large neglected part of the parameter space in pulsar studies is that of observing length -- quite naturally as many time allocation committees do not appreciate long-term projects. However, some new phenomena will only be discovered, when sources -- not only pulsars, that is --- are monitored for a long time span. In this review I will concentrate on the period space, and the implication that can be derived from studying emission of millisecond pulsars. For the other aspects, the reader may be referred to other contributions (e.g.~Karastergiou et al.) or recent literature (e.g.~Kramer et al.~2002). ",
        "conclusions": "The emission properties of millisecond pulsars are in many respects similar to those of slowly rotating pulsars. However, there are a few remarkable differences like additional profile components and a very week frequency evolution for most MSPs, which can be attributed to the smaller, compact magnetosphere. The additional pulse components may be representatives of outer gap emission, providing an interesting link to the X-ray properties of MSPs. This motivates a close inspection of their joint radio and high energy characteristics. Results of this work in progress will be presented elsewhere."
    },
    "0206/hep-ph0206097_arXiv.txt": {
        "abstract": "In the mean field approximation, we study the effects of weak magnetism and pseudoscalar interaction in the neutrino energy losses caused by the direct Urca processes on relativistic nucleons in the degenerate baryon matter. Our formula for the neutrino energy losses incorporates the effects of nucleon recoil, parity violation, weak magnetism, and pseudoscalar interaction. For numerical testing of our formula, we use a self-consistent relativistic model of the multicomponent baryon matter. We found that, due to weak magnetism effects, relativistic emissivities increase by approximately 40-50\\%, while the pseudoscalar interaction only slightly suppresses the energy losses, approximately by 5\\%. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206325_arXiv.txt": {
        "abstract": "In this paper we present the mass functions in the substellar regime of three young open clusters, IC 348, $\\sigma$ Orionis and Pleiades, as derived using the data from the 2 Micron All Sky Survey (2MASS) catalogue which has a limiting magnitude of $K_{s} \\sim 15$, and the latest version of the Guide Star Catalogue (GSC) which has a limiting magnitude of \\setcounter{footnote}{3} $F\\footnote{F refers to the POSS II IIIa-F passband} \\sim$ 21. Based on recent evolutionary models for low mass stars, we have formulated the selection criteria for stars with masses below $0.5 M_{\\odot}$. Using a statistical approach to correct for the background contamination, we derive the mass function of objects with masses ranging from 0.5$ M_{\\odot}$ down to the substellar domain, well below the Hydrogen Burning Mass Limit.  The lowest mass bins in our analysis are 0.025, 0.045 and 0.055 M$_\\odot$ for IC 348, $\\sigma$ Orionis and Pleiades, respectively. The resultant slopes of the mass functions are 0.8 $\\pm$ 0.2 for IC 348,  1.2 $\\pm$ 0.2 for $\\sigma$ Orionis and 0.5 $\\pm$ 0.2 for Pleiades, which are consistent with the previous results. The contribution of objects below 0.5 M$_\\odot$ to the total mass of the cluster is $\\sim$40\\%, and the contribution of objects below 0.08 M$_\\odot$  to the total mass is $\\sim$4\\%. ",
        "introduction": "The Initial Mass Function (IMF) of stars is one of the most fundamental and crucial ingredients in models of galaxy formation and stellar evolution. It determines several key parameters in stellar populations, such as the yield of heavy elements, the mass-to-light ratio, luminosity evolution over time, and the energy input into the interstellar medium. The determination of the IMF is therefore of great astrophysical importance. The IMF of low-mass stars  is of special interest in this context, since they contain a major  fraction of the stellar mass, and have been hypothesized to contain a significant fraction of  the total mass in the Universe (see, e.g., Fukugita, Hogan and Peebles, 1998).  In this paper we mainly deal with objects having masses less than  0.5$M_{\\odot}$, the lowest mass of the detectable objects being as low as 0.025$M_{\\odot}$. These low-mass stars evolve little over the lifetime of the Universe, and hence the observed present day Mass Function of these stars is likely to be a good representation of their IMF. But, the IMF at or below the HBML remains poorly known,  mainly for two reasons. First, such objects are faint and hence difficult to detect. Second, the mass-luminosity relation of these objects is uncertain and model-dependent, and hence their mass determination is imprecise. Significant improvements are being made on both these aspects, as described below.  The difficulty caused by their faintness can be greatly alleviated by concentrating on young, low-mass objects since the low-mass stars at or below the HBML are expected to be warmer and more luminous when young, although they rapidly cool and fade with age (Burrows et al. 1997, D'Antona \\& Mazzitelli 1997, Baraffe et al. 1998). Hence young and nearby open clusters provide a good opportunity to study the low end of the stellar IMF, since the suitable combination of their youth and proximity makes it possible to detect objects well below the  HBML in these clusters, particularly at near-infrared wavelengths.  The uncertainty in the mass-luminosity relation has been greatly reduced by the tremendous progress in the theoretical models for the evolution of these cool and dense objects over the past few years. These models play a crucial role in predicting masses of these low-mass objects from the observable quantities like colors and luminosities. Burrows et al. (1997) have generated models of spectra, colors and evolution of brown dwarfs using nongrey calculations. Their models span the mass range of 0.3$M_{J}$ to 70$M_{J}$ (where $M_{J}$ refers to the mass of Jupiter) with effective temperatures varying from $\\sim$1300 K to 100 K. D'Antona \\& Mazzitelli (1997) simulate the evolution of objects in the mass range 20$M_{J} \\leq M \\leq 1.5 M_{\\odot}$. They describe the star's evolution from the hydrostatic phases of pre-main sequence contraction to the hydrogen burning main sequence phase through deuterium and lithium burning . Baraffe et al.(1998) have developed the evolutionary models in the 0.075 $M_{\\odot}$ to 1 $M_{\\odot}$ mass range for solar type metallicities based on the NextGen atmospheric models of Allard et al. (1996), and Chabrier et al. (2000) have extended this study by including dust formation and opacity. The recent models for dwarfs by Marley et al (2002)  take the extra effect of sedimentation into account, which suggest that some of the colors, particularly the Sloan i$'$-z$'$, can be greatly affected by sedimentation.  The advent of the red-sensitive CCDs and 2-dimensional near-IR detectors in the last decade has made it possible to detect such low-mass objects, and there have been numerous imaging surveys targeted towards open clusters to probe the substellar domain. Surveys by Wilking et al. (1999) and Luhman et al. (1999) for $\\rho$ Ophiuchi, Herbig (1998) and Luhman (1999) for IC 348, B\\'{e}jar et al. (1999) and Zapatero Osorio et al. (1999a) for $\\sigma$ Orionis, Zapatero Osorio et al. (1996) and Stauffer et al. (1999) for $\\alpha$ Persei, Zapatero Osorio et al. (1997;1999b), Bouvier et al. (1998) and Hambly et al. (1999) for Pleiades, Hambly et al. (1995), Pinfield et al. (1997) and Magazz\\`{u} et al.(1998) for Praesepe, Gizis et al. (1999) and Reid \\& Hawley (1999) for Hyades, and Barrado y Navascu\\'{e}s et al. (2001a) for IC 2391 are to name a few.  The recent release of the 2MASS catalogue in the near-infrared wavelengths with a limiting magnitude of $K_{s} \\sim 15$, and the latest (development) version of the Guide Star Catalogue (GSC) with a limiting magnitude of $F \\sim 21$ form an ideal combination to study low-mass objects in nearby open clusters using a statistical approach. We have used these two datasets in conjunction with the recent evolutionary models to isolate the low mass members of the clusters. The cluster membership is not ascertained by follow-up spectroscopy or proper motion studies. But the background/foreground contamination is accounted for statistically by studying nearby control fields.  In \\S 2, we review some previous work on the derivation of the mass function; in \\S 3, we describe the rationale of our sample selection and the details of the 3 individual clusters selected for this study; in \\S 4, we describe the data, the procedure  adapted in selecting the cluster members and the method used for their mass determination; in \\S 5, we describe the specific selection criteria used for the  individual clusters and the slope of the resultant mass function; and  we end with a discussion of the results in \\S 6. We plan to extend this work to more clusters in the future. ",
        "conclusions": "Recent surveys have found a significant population of low mass stars, brown dwarfs and planetary mass objects in young open clusters. We have adopted a statistical approach to determine the mass spectrum, $dN/dM \\propto M^{-\\alpha}$, of objects in the mass range 0.5$M_{\\odot}$ to 0.025-0.055$M_{\\odot}$, using the data from the recently released 2MASS and the GSC catalogues. Unlike some of the previous studies, our study makes use of both the optical data as well as the near-IR data. Since these datasets cover a large portion of the sky, they allow us to study the entire area covered by each cluster. As a result, the areas covered in our study are generally larger than the areas covered in most of the previous studies. These datasets also allow us to apply a statistical approach to efficiently subtract the background contribution using several control fields close to the cluster.  We carried out a detailed study for IC 348. Fo this cluster, we derived the mass functions using the solar metalicity model of Baraffe et al. (1998) and the dusty models of Chabrier et al. (2000), both of which gave very similar results. We also compared the confirmed low-mass members from Luhman (1999) with our isolated low-mass candidates which further strengthened the validity of this technique. We then used the same technique to $\\sigma$ Orionis and Pleiades. The resultant slopes of the mass functions for IC 348, $\\sigma$ Orionis and Pleiades are 0.7, 1.2 and 0.5 respectively, with an estimated error of $\\pm$0.2. For IC 348 we have used the mass range from 0.5$M_{\\odot}$ to 0.025$M_{\\odot}$ in deriving the slope of the mass function whereas for $sigma$ Orionis and Pleiades, the lowest mass bins correspond to 0.045 $M_{\\odot}$ and 0.055$M_{\\odot}$, respectively. Taking into consideration the effect of mass segregation and preferential loss of low mass members from the cluster, the mass function derived here for the inner 1.5 deg radius of Pleiades cluster could be considered as a lower limit to the true mass function. As discussed under individual sections, the mass functions derived here are in good agreement with the value derived by other groups based on studies of confirmed low mass members of these clusters, which demonstrates the consistency of different approaches. Within the uncertainties, the values of $\\alpha$ derived for the IC 348 and $\\sigma$ Orionis is in agreement with that derived for the local sample of low-mass stars (Reid et al. 1999). The derived slopes imply that the mass spectrum continues to rise well below the HBML, but the mass functions are appreciably flatter in the low-mass regime than the Salpeter mass function. The results are summarized in Table 3.     Taking the Salpeter exponent of 2.35 in the mass range 1 -- 10 $M_{\\odot}$, the Chabrier exponent of 1.55 in the mass range 0.5 -- 1 $M_{\\odot}$, and the values obtained by us below 0.5, we calculate the mass contribution to be about 40\\% for objects below 0.5 M$_\\odot$, and about 4\\% for objects below the HBML of 0.08 M$_\\odot$. (Note that the contributions are not sensitive to the choice of the slope in the higher mass regime. For example, if we use the value of $\\alpha$ as 2.7 for M$> 1 M_\\odot$ as derived by Chabrier (2001) instead of the Salpeter value of 2.35, the mass contributions change only by $\\sim1\\%$). Our results are consistent with that of the previous studies (e.g. B\\'{e}jar et al. 2001), and suggest that, although the low mass stars are at least as numerous as their high mass counter parts (as seen from figures 5,7 and 9), their contribution to the total mass is small. The contributions of low-mass objects to the total mass in the clusters seem to be marginally smaller than that of the low-mass objects in the local sample (e.g. Reid et al. 1999), but the slope of the mass function is less steep for the relatively older Pleiades cluster. This is not surprising since the high-mass stars are likely to be preferentially lost in an older and mixed population such as the local sample.  Follow up spectroscopic observations to confirm the isolated low mass members of the clusters would further strengthen the results derived from this purely statistical approach."
    },
    "0206/astro-ph0206055_arXiv.txt": {
        "abstract": "{ We present new Near-Infrared (NIR) observations, in the J, H and Kn bands, for a sample of Polar Ring Galaxies (PRGs), selected from the Polar Ring Catalogue (Whitmore et al. \\cite{PRC}). Data were acquired with the CASPIR near-IR camera at the 2.3 m telescope of Mount Stromlo and Siding Spring Observatory. We report here on the detail morphological study for the central host galaxy and the polar structure in all PRGs of our sample. Total magnitudes, bulge-to-disk decomposition and structural parameters are computed for all objects. These data are crucial for an accurate modeling of the stellar population and the estimate of the star formation rates in the two components. ",
        "introduction": "Polar Ring Galaxies (PRGs) are classified as dynamically peculiar systems, as they show the coexistence of two luminous components, the central host  galaxy and ring, with their angular momentum vectors in two nearly orthogonal  planes (Schweizer et al. \\cite{Schweizer83}; Whitmore et al. \\cite{PRC}). The presence of almost two perpendicular angular momentum vectors cannot be explained through the collapse of a single protogalactic cloud: a ``second event'' must have occurred in the formation history of these systems. In the last years, a number of observational studies were produced to constrain the origin of PRGs (Reshetnikov \\cite{Resh97}; see also the review by Sparke \\& Cox \\cite{Cox2000}). In almost all PRGs the morphology of the host galaxy resembles that of an early-type object (elliptical or S0 galaxy), because of its structure-less appearance and no HI: kinematical studies on some PRGs have confirmed that this component is rapidly rotating (Schechter et al. \\cite{Schechter84}; Whitmore et al. \\cite{PRC}). The integrated colors and gas-to-dust ratio, together with the large $M(HI)/L_B$ ratio for the whole system suggest that PRGs may be quite similar to the late-type spirals (Arnaboldi et al. \\cite{magda95}). Very recent works on NGC~4650A, which is considered the prototype of the class of wide PRGs, (Arnaboldi et al. \\cite{magda97}; Iodice et al. \\cite{4650aI}; Gallagher et al. \\cite{4650aG}) have shown that the polar structure appears to be a disk of a very young age; moreover the host galaxy integrated colors and light distribution do not resemble that of a typical early-type system.  The main goal of the present work (Paper I) is to provide accurate photometry in the NIR for a sample of PRGs; in a second paper (Iodice et al. \\cite{paperII}, Paper II) they will be compared with the predictions from different formation scenarios for these peculiar systems. Near-IR photometry is necessary to reduce as much as possible the dust absorption that strongly affects the starlight distribution in the host galaxy and in the ring (Whitmore et al. \\cite{PRC}). Since the dust optical depth decreases toward longer wavelengths, photometry in the NIR will be relatively free from this problem, and the inner structures of the host galaxy and ring may be easily identified. The NIR photometry is also more representative of the light emitted by the older stellar population, which contains most of the mass: any dynamical modeling of PRGs will be more accurate when using the light distribution for the different components in the NIR rather than those at the optical wavelength. In addition, the study of optical and NIR integrated colors will yield information about the age and metallicity of dominant stellar population in the different components of a PRG system.  In this paper we present new NIR observations, obtained for a selected sample of PRGs from the Polar Ring Catalogue, listed in Table~\\ref{prg}, and have applied the same procedures adopted to study the polar ring galaxy NGC~4650A (Iodice et al. \\cite{4650aI}). Observations and data reduction are presented in Sec.\\ref{observations}; the morphology, light and color distribution of the two components (host galaxy and ring) are discussed in Sec.\\ref{morphology} and Sec.\\ref{photometry}. In Sec.\\ref{color} the integrated colors derived for different regions of each PRG are presented. The two-dimensional model of the host galaxy light distribution is discussed in Sec.\\ref{2Dmodel},a detailed description of each selected PRG is given in Sec.\\ref{obj_descr}. The final summary of the data is presented in Sec.\\ref{conclu}.  \\begin{table*} \\centering \\caption[]{The Polar Ring Galaxy sample studied in this work. In the second column of the table we list the  object identification as given in the Polar Ring Catalogue, PRC, (Whitmore et al. \\cite{PRC}); coordinates $\\alpha$ and $\\delta$, the Heliocentric velocities, and galaxy extension (derived from NED database) are reported in the third, fourth, fifth and sixth columns respectively.} \\label{prg} \\begin{tabular}{cccccc} \\hline\\hline Object name & PRC name & $\\alpha$ (J2000) & $\\delta$ (J2000) & $V_{0}$ & diameters\\\\ &  &  &  & (km/s) & (arcmin)\\\\ \\hline A0136-0801 & A-01 & 01h38m55.2s & -07d45m56s & 5500 & 0.41 x 0.3\\\\ ESO 415-G26 & A-02 & 02h28m20.1s & -31d52m51s & 4604 & 1.3 x 0.6 \\\\ ARP 230 & B-01 & 00h46m24.2s & -13d26m32s & 1742 & 1.3 x 1.2\\\\ AM 2020-504 & B-19 & 20h23m54.8s & -50d39m05s & 4963 &   \\\\ ESO 603-G21 & B-21 & 22h51m22.0s & -20d14m51s & 3124 & 1.1 x 0.6 \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{obj_descr} Here we review in detail the main features which characterize the light and color distribution in the NIR for each PRG in our sample. We will also review other important properties of these objects reported in literature, which we will use in Paper II. \\bigskip  {\\bf A0136-0801} - This is one of the best case of kinematically confirmed polar ring galaxy (Schweizer et al. \\cite{Schweizer83}; Whitmore et al. \\cite{PRC}, PRC): it is characterized by a wide polar structure, which is three times more extended than the optical radius of the central host galaxy. Along the polar ring major axis the surface brightness profile in the H band (see Fig.\\ref{prof}) is given by two exponential segments with different slopes: the light inside $10''$ from the center is associated with the host galaxy, whereas the ring component is extended out to $40''$. Both the H band image and luminosity profiles (Fig.\\ref{hfm1} and Fig.\\ref{prof}) show that the polar ring is less luminous than the central galaxy, and the bulk of H band light is concentrated at smaller radii. The high-frequency residual image in the H band (see Fig.\\ref{hfm1}) shows the presence of a nearly edge-on disk in the central host galaxy. From the 2MASS database we have derived the J, H and Ks data for this object and we have computed the total integrated magnitudes in a circular area, with a $40''$ radius in order to include the whole ring. This is the same area used to compute the total magnitude in CASPIR H band (see Tab.~\\ref{totm}). Due to the very short exposure (1.3 sec), in the 2MASS images the polar structure is very faint, and most of the light comes from the central galaxy. Thus the J-H and H-K colors derived from these data ($J-H=0.41$, $H-Ks=0.43$) are more rapresentative of the host galaxy, and they are comparable with the J-H and H-K colors obtained for host galaxy of other PRGs in our sample. This object was mapped in HI with the {\\it Very Large Array (VLA)} by van Gorkom et al. (\\cite{vgorkom87}) and Cox (\\cite{Cox96}): all HI emission is found within the polar ring, whose outer HI contours appear to warp away from the poles. The total HI mass estimated for this object is about $1.6 \\times 10^9 M_{\\odot}$ (Cox \\cite{Cox96}). The regular HI distribution and optical appearance, an apparent lack of HII regions and  other signs of recent star-formation activity (Mould et al. \\cite{Mould82}) suggest that the polar structure is quite old and possibly dynamically stable.  \\bigskip  {\\bf ESO 415-G26} - This is a well-known polar ring galaxy. Unlike A0136-0801, the polar ring is less extended than the central host galaxy in the optical band (see Fig.\\ref{eso415_B}). Deeper exposures show a lot of debris at a position angle which is intermediate between those of the host galaxy major axis and the ring. They also show shells and loops in the outer regions (Whitmore et al. \\cite{whitmore87}). In the NIR bands (Fig.\\ref{hfm1}, middle panels) the polar ring is so faint that it is hardly detected. The central host galaxy is the dominant luminous component: the analysis of the light distribution shows that this is a nearly edge-on S0 galaxy with an exponential bulge. Among the PRGs of our sample, ESO~415-G26 is characterized by a highest B/D ratio (see Table~\\ref{2dparam}). The NIR color maps (Fig.\\ref{mappe}) and J-H vs. H-K colors in different regions of host galaxy and ring  (see Sec.\\ref{color}) show that the nucleus of the system is characterized by the reddest color and that the polar ring is much bluer than the central host galaxy. The HI map for this object was obtained with the VLA by van Gorkom et al. \\cite{vgorkom87} and by van Gorkom \\& Schiminovich \\cite{vgorkom97}. They noted that the neutral hydrogen lies along the major axis of the polar ring, with some degrees of correlation between the HI and the outer shells (van Gorkom \\& Schiminovich \\cite{vgorkom97}). The most accurate estimate of the total HI mass is about $5.6 \\times 10^9 M_{\\odot}$ (Schiminovich et al. \\cite{Schiminovich97}). This object is characterized by a considerable amount of molecular hydrogen: Galletta et al. (\\cite{galletta97}) estimated the total $H_{2}$ mass to be about $2.4 \\times 10^9 M_{\\odot}$.  \\bigskip  {\\bf ARP 230} - This object, also known as IC 51, was studied by Wilkinson et al. (\\cite{wilkinson87}) as a well-known example of {\\it shell elliptical} galaxies: in the NE and SW directions, outer shells are clearly visible (Hernquist and Quinn \\cite{Hernquist88}), which are more luminous in the B band (see Whitmore et al. \\cite{PRC}, PRC) than in the NIR ones (Fig.\\ref{hfm1},  bottom panels). It is also classified as a PRG because it has a fast rotating disk-like structure, made up by gas, stars and dust, perpendicular to the apparent major axis of the central galaxy, $P.A.=37^{\\circ}$,  (Mollenhoff et al. \\cite{Mollenhoff92}). Both optical and NIR images (Fig.\\ref{hfm1}, bottom panels) of this polar ring  galaxy show that the ring-like component, along SE and NW directions, has the size of the inner galaxy, and it has a very well-defined outer edge, where dust absorption is present. The high-frequency residual images (Fig.\\ref{hfm1} right panels), show a very distorted  structure for the ring component: it seems strongly warped at about $10''$ from the center, with associated absorption features. An elongated structure nearly orthogonal to the ring suggests that the central host galaxy is more similar to a disk galaxy, an S0, than an elliptical galaxy. In this case it is very difficult to distinguish the morphology of the central component. The 2D model of the light distribution also suggests that it may be an S0 galaxy with an exponential bulge (Table~\\ref{2dparam}), but the peculiar ring structure produces a very uncertain estimate of all the structural parameters; the apparent axial  ratios at larger radii are also influenced by the presence of the outer shells. The bright edges of the ring (in the NW and SE directions) and outer shells are clearly visible as bright residuals in the image ratio between the whole galaxy and the 2D model of the central component (see Fig.\\ref{hfm1}, lower left panels). The central host galaxy shows similar J-H colors to the polar ring component, and bluer H-K colors (see Table~\\ref{polycol}). The very red colors of the nucleus of the system are possibly due to dust absorption in the ring, which passes in front it. Near-IR images, in J, H and Ks bands, are also available for this object in the 2MASS database. Taking into account the average uncertainties which affect these data (see Sec.\\ref{color}), the 2MASS total magnitudes are comparable with those derived from the CASPIR data (see Tab.\\ref{totm}. This object was mapped in HI of the VLA by Schiminovich et al. (\\cite{Schiminovich97}),  and they found that the neutral hydrogen is all associated to the ring and shows rotation along this component. They estimated a total HI mass of about $2.3 \\times 10^9 M_{\\odot}$. This PRG was also mapped by Cox (\\cite{Cox96}) in the radio continuum, at 20 cm and 6 cm, with the VLA: she found an extended emission aligned with the ring structure and additional filaments which are extended above the ring plane. By comparing the radio continuum and the far-infrared (FIR) emission\\footnote{The far-infrared (FIR) emission was detected $60\\mu m$ and $100\\mu m$ by Moshir et al. \\cite{Moshir90}, the IRAS Faint-Source Catalog)}, Cox (\\cite{Cox96}) deduced that this PRG falls on the radio/FIR correlation for star-forming galaxies.  \\bigskip  {\\bf AM 2020-504} - Previous photometric and spectroscopic observations showed that the central host galaxy in this object is very similar to an elliptical galaxy, which is characterized by a decoupled rapidly rotating core within $3''$ from the center, (Whitmore et al. \\cite{whitmore87}; Arnaboldi et al. \\cite{magda93}, \\cite{magda95}). The narrow polar ring, which is observed along the host galaxy minor axis, is brighter in the optical than in the NIR images (Fig.\\ref{hfm1}). In all bands, the light distribution of this component is peaked between $10''$ and $15''$ (see Fig.\\ref{prof}). As pointed out in the previous Sections, this object shows different properties with respect to the other polar ring galaxies studied here. In the high-frequency residual images (Fig.\\ref{hfm1}, middle panel), there is no trace of any disk-like structures associated with the host galaxy major axis,  which is observed in all the other polar ring galaxies studied here. The absence of a disk in the host galaxy suggested the use of a Sersic law (Eq.\\ref{sersic}) for the 2D fit of the light distribution in this component; the results indicate that the AM 2020-504 central component is an elliptical rather than an S0 galaxy. \\bigskip  {\\bf ESO 603-G21} - The prominent structure which appears in the B band image of this object is the warped dusty ring (in the SE and NW directions) which surrounds a bright round stellar system (see Whitmore et al. \\cite{PRC} and Arnaboldi et al. \\cite{magda95}). This central component is much fainter in the NIR images (Fig.\\ref{hfm1}) and is embedded in a very luminous disk-like structure. The high-frequency residual images (Fig.\\ref{hfm1}) reveal that this disk is nearly edge-on. A further fainter ``filamentary'' structure, which was already detected by Arnaboldi et al. (\\cite{magda95}), is visible perpendicular to this disk, and aligned with the apparent major axis of the central spheroid ($P.A.=24^o$). The extension of this filamentary structure is less than $10''$. The reddest regions in the NIR color maps (Fig.\\ref{mappe}) correspond to the disk component, while the central spheroid is bluer. ESO~603-G21 is the object in our sample with the reddest nuclear regions (see Table~\\ref{polycol}). Very recently, Reshetnikov et al. (\\cite{Resh02}) have performed a detailed surface photometry of ESO~603-G21 in the optical B, V and R bands: they found that the central component has an exponential light distribution and is surrounded by an extended, warped, edge-on disk/ring structure. These results are consistent with our findings.\\\\ The surface brightness profiles along the bright edge-on disk have an exponential behavior (see Fig.\\ref{prof}) and the comparison with the surface brightness profiles along the orthogonal direction shows that this disk is the dominant luminous component also in the NIR bands, unlike all the other polar rings which are faintest in the NIR. This peculiarity makes ESO~603-G26 similar to a late-type spiral galaxy with a kinematically-decoupled extended bulge (e.g. NGC~4672 and NGC~4698, see Bertola et al. \\cite{bertola99} and Sarzi et al. \\cite{sarzi2000}) rather than a polar ring galaxy, as also suggested by Reshetnikov et al. (\\cite{Resh02}).  \\\\ As we will show in Paper~II, the central component of ESO~603-G26 has colors, age and light distribution properties quite similar to those of the host galaxy in the other PRGs of our sample. In particular, the surface brightness distributions is well-fit by the super-position of two components (bulge + disk) rather than by only an $r^{1/4}$ bulge, as in NGC~4672 and NGC~4698. The available spectroscopic data for this object, which should help us to understand what kind of object ESO~603-G26 is, are quite uncertain (Arnaboldi et al. \\cite{magda95}): they show rotation of the stellar component along the two axis corresponding to $P.A.=24^\\circ$ and $P.A.=114^\\circ$. This result may suggest that the underlying central spheroid is triaxial. The rotation curve derived from the strong $H\\alpha$ emission seen in the disk/ring spectrum shows a constant velocity gradient along this component (Arnaboldi et al. \\cite{magda95}). This may suggest that we are indeed looking at a ring. However, these data are strongly influenced by the dust absorption, so no definite conclusion can be derived.\\\\ The J, H and Ks total magnitudes derived from the 2MASS data, available for this system, are on average 0.16 mag brighter than the J, H and Kn CASPIR magnitudes (see Tab.\\ref{totm}). As already discussed in Sec.\\ref{totm}, the non-Poissonian background fluctuations, which affects both CASPIR and 2MASS images, particularly in H band, can explain such differences.\\\\ Radio data for this object shows CO emission corresponding to $1.1 \\times 10^9 M_{\\odot}$ of molecular hydrogen (Galletta et al. \\cite{galletta97}) and $6.2 \\times 10^9 M_{\\odot}$ in HI (van Driel et al. \\cite{vdriel2000}). Radio continuum emission was detected for this object by Cox (\\cite{Cox96}) in the central regions of the candidate polar ring (i.e. along $P.A.=114^o$, as in the PRC). FIR emission, at $60\\mu m$ and $100\\mu m$, was detected for this PRG (Moshir et al. \\cite{Moshir90}, the IRAS Faint-Source Catalog), this object too falls on the radio/FIR correlation for star-forming galaxies rather than for an AGN (Cox \\cite{Cox96})."
    },
    "0206/astro-ph0206263_arXiv.txt": {
        "abstract": "We present 3 years of photometry of the ``Double Hamburger'' lensed quasar, HE1104$-$1805, obtained on 102 separate nights using the OGLE 1.3-m telescope.  Both the A and B images show variations, but with substantial differences in the lightcurves at all time delays.  At the $310^{\\rm d}$ delay reported by Wisotzki and collaborators the difference lightcurve has an rms amplitude of 0.060 mag.  The structure functions for the A and B images are quite different, with image A more than twice as variable as image B (a factor of 4 in structure function) on timescales of less than a month.  Adopting microlensing as a working hypothesis for the uncorrelated variability, the short timescale argues for the relativistic motion of one or more components of the source.  We argue that the small amplitude of the fluctuations is due to the finite size of the source with respect to the microlenses. ",
        "introduction": "Two very different physical processes contribute to the observed photometric variability of gravitationally lensed quasars: the intrinsic variabilty of the quasar itself and propagation effects along the line of sight.  Chief among the latter is microlensing by the stellar mass objects in the intervening lens (Chang and Refsdal 1979; Paczy{\\'n}ski 1986).  The combination of intrinsic and microlensing variations represents an embarassment of riches.  For the purpose of measuring lens time delays (using the correlated intrinsic variabilty of the quasar images) uncorrelated microlensing variations are an additional source of noise.  Conversely, the intrinsic variation of the quasar produces correlated noise in the uncorrelated microlensing signals.  Time delays have been measured for nearly a dozen systems, and in most cases microlensing appears not to have presented a serious problem (e.g. Kundi{\\'c} et al. 1997; Schechter et al. 1997).  Dramatic microlensing variations have been observed in the system 2237+0305 (Corrigan et al. 1991; Wo{\\'z}niak et al. 2000) but on a timescale (months) which is very much longer than the predicted delays (hours).  There have, however, been instances in which microlensing and time delay measurements have interfered with each other.  In the case of 0957+561, the two images have long timescale (1000$^{\\rm d}$) variations (Refsdal et al. 2000), which bias the inferred time delay.  Burud et al. (2000) report uncorrelated variations over timescales of several months in their study of B1600+434.  Examination of their Figure 3 shows apparent uncorrelated variations on timescales of days. Uncorrelated variations are also reported by Hjorth et al.\\ (2002) in their study of RXJ0911+0551.  We report here the results of an unsuccessful program to measure the time delay of the doubly imaged quasar HE1104$-$1805 (Wisotzki et al. 1993).  In three years' monitoring with the Optical Gravitational Lensing Experiment (OGLE) 1.3m telescope at Las Campanas we see uncorrelated variations in the A and B images, which we interpret as the result of microlensing.  In \\S 2 we describe the observations and initial reductions.  In \\S 3 we compare light curves for the two quasar images, A and B, using the time tested chi-by-eye technique and a less subjective method.  Our data fail to produce a satisfactory time delay.  In \\S 4 we derive structure functions separately for the A and B images.  Adopting the time delay measured by Gil-Merino et al.\\ (2002), we determine a structure function for the microlensing from the difference between the A and B images.  In \\S 5 we explore several alternative interpretations of the structure functions for the two quasar images. ",
        "conclusions": "Three years of observations of the two lensed quasar images of HE1104$-$1805 show variations in the A and B images that are uncorrelated, with V amplitudes of $\\sim 0.060$ mag. The A image exhibits considerably more variability, on a timescale $\\lesssim$ 1 month, while the fluctuations in the B image are consistent with our expectations for variations intrinsic to the quasar.  On the hypothesis that the fluctuations are due to microlensing by solar mass stars, the implied source velocity is $\\sim 0.25 c$.  For reasonable assumptions regarding the ratio of dark to microlensing matter and on the assumption that only a single hotspot is contributing to the fluctuations, the hotspot contributes $\\sim 7$\\% of the continuum flux.  A multiple hotspot model presented (and rejected) by Wyithe and Loeb (2002) also seems viable, perhaps even preferable, in the present case."
    },
    "0206/astro-ph0206505_arXiv.txt": {
        "abstract": "The most energetic particles ever detected exceed $10^{20}$ eV in energy. Their existence represents at the same time a great challenge for particle physics and astrophysics, and a great promise of providing us for a probe of the validity of the laws of Nature in extreme conditions. We review here the most recent data and the future perspectives for detection of cosmic rays at ultra-high energies, and discuss possible ways of using these data to test the possibility that new Physics and/or new Astrophsyics may be awaiting around the corner.\\par\\vskip .3cm  \\hskip 6cm {\\footnotesize \\it ``When you carry out an experiment there are two possible outcomes:}\\par \\hskip 6cm {\\footnotesize\\it either you confirm the theoretical expectation, and in this case you made}\\par \\hskip 6cm {\\footnotesize \\it a measurement, or you don't, and in this case you made a discovery''.}\\par \\hskip 14cm {\\small\\it E. Fermi} \\vspace{1pc} ",
        "introduction": "The cosmic ray spectrum has been now measured over a large range of energies, that extends in its upper part to more than $10^{20}$ eV, the so-called {\\it ultra high energy cosmic rays} (UHECRs). The quest for the origin of these high energy particles over the whole range is still open and represents one of the big challenges for the future. In the highest energy end of the cosmic ray spectrum, several issues make the challenge even harder: a) it is hard to envision possible acceleration sites where particles with energy in excess of $10^{20}$ eV may be accelerated; b) even if some class of sources could indeed accelerate particles to the highest energies, a homogeneous spatial distribution of these sources would leave an imprint in the cosmic ray spetrum, known as Greisen-Zatsepin-Kuzmin (GZK) cutoff, due to photopion production on the photons of the cosmic microwave background \\cite{greisenzk66}. On the basis of the first issue, we would expect an {\\it end of the cosmic ray spectrum} to occur at some {\\it high} energy. On the basis of the second issue, we would expect a strong flux suppression (not really a cutoff) at about $5\\times 10^{19}$ eV.  Several experiments have been operating to detect the flux of UHECRs, starting with Volcano Ranch \\cite{linsley} and continuing with Haverah Park \\cite{watson91} and Yakutsk \\cite{efimov91} to the more recent experiments like AGASA \\cite{agasa01,tak99,tak98,ha94}, Fly's Eye \\cite{bird93,bird94,bird95} and HiRes \\cite{kieda99}.  At present there is no clear indication that the cosmic ray spectrum comes to an end due to either one of these two reasons. More statistics of events is however required to achieve a solid conclusion in this respect.  The physics involved in the explanation of the origin and propagation of UHECRs needs often to be pushed to its extremes to accomodate observations. This transforms a problem into a precious tool to probe a territory which is still uncharted, a New Physics which several hints tell us should exist, but is still hidden somewhere. In the following we describe a few directions in which this investigation may lead.  In \\S \\ref{sec:observations} we briefly summarize the status of current observations; in \\S \\ref{sec:gzk} we restate the basic issues that are known as GZK problem; in \\S \\ref{sec:NP} we illustrate two examples of new physics that can be investigated with the help of UHECRs, namely Physics close to the grand unification scale (\\S \\ref{sec:td}) and possible violations of Lorentz invariance (\\S \\ref{sec:li}). In \\S \\ref{sec:conc} we give our conclusions. ",
        "conclusions": "}  The search for the end of the cosmic ray spectrum, started a few decades back in time, is still ongoing and still not successful. This challenge led us to the detection of particles with energies in excess of $10^{20}$ eV. Acceleration processes are strongly limited by energy losses and finite size of the known acceleration regions and only some types of sources are barely able to energize protons up to the observed energies \\cite{olinto}. Composition and anisotropy studies will be the keys to solve the mystery, but at the cost of increasing the statistics by at least a factor of 10 compared with current experiments. Two experiments are being planned for the next decade or so, and will provide the characteristics necessary to do cosmic ray astronomy: the Auger project \\cite{cronin} is currently in the construction stage in Argentina, while the EUSO project \\cite{scarsi} is scheduled for operation starting in 2008. Each one of these enterprises implies an improvement by a factor $\\sim 10$ compared with the previous one, which means a predicted 500 events per year above $10^{20}$ eV for EUSO, if the AGASA spectrum is taken as a template.  Besides being the tools for ultra high energy cosmic ray astronomy, these experiments represent a unique tool to study possible New Physics at extremely high energies. The case of neutrino oscillations provides an example of the first hint of the existence of Physics beyond the Standard Model of Particle interactions, derived in an Astrophysics context. It is foreseeable that the ball of particle physics, after a few decades, could go back to the field of cosmic rays, where the first steps in that direction were moved in the '30s."
    },
    "0206/astro-ph0206219_arXiv.txt": {
        "abstract": "A new, complete, theoretical rotational and vibrational line list for the $A~^2\\Pi \\leftarrow X~^2\\Sigma^+$ electronic transition in MgH is presented. The list includes transition energies and oscillator strengths for all possible allowed transitions and was computed using the best available theoretical potential energies and dipole transition moment function with the former adjusted to account for experimental data. The $A\\leftarrow X$ line list, as well as new line lists for the $B'~^2\\Sigma^+ \\leftarrow X~^2\\Sigma^+$ and the $X~^2\\Sigma^+ \\leftarrow X~^2\\Sigma^+$ (pure rovibrational) transitions, were included in comprehensive stellar atmosphere models for M, L, and T dwarfs and solar-type stars. The resulting spectra, when compared to models lacking MgH, show that MgH provides significant opacity in the visible between 4400 and 5600\\AA~. Further, comparison of the spectra obtained with the current line list to spectra obtained using the line list constructed by \\citet{kur93} show that the Kurucz list significantly overestimates the opacity due to MgH particularly for the bands near 5150 and 4800\\AA~ with the discrepancy increasing with decreasing effective temperature. ",
        "introduction": "The study of the spectra of cool stars requires detailed knowledge of molecular opacities. This includes important absorbers such as TiO, CO, and water vapor, which have bands that cover large wavelength ranges and are very important for the structure of the atmosphere due to their overall cooling or heating effects.  In addition, there are a number of molecules that have bands covering comparatively small wavelength ranges (e.g., a few 10 or $100\\,$\\AA). Many of them are trace molecules that have only small effects on the overall physical conditions inside the atmosphere but that are important for spectral classification and for the determination of stellar parameters such as effective temperatures, gravities and abundances. Unfortunately, important molecular data such as energy levels, bound-free, and bound-bound cross-sections are only poorly known or not known at all for a number of these trace molecules. We have therefore started a project to update or provide for the first time molecular data of astrophysical interest for important trace molecules and consider in this work MgH. These data will be computed using state-of-the-art molecular physics codes and should improve our ability to model and analyze cool stellar atmospheres considerably.  It is important to assess the quality of the computed molecular data. This is best done by comparing to experimental results; however, this is only possible for very few molecules of astrophysical interest. In addition, in many cases the temperature range of astrophysical importance is higher than what can be reached with current experimental setups. Therefore, indirect methods of testing and evaluating the molecular data are useful. In this paper, we use the general-purpose stellar atmosphere code {\\tt PHOENIX} to calculate model atmospheres and synthetic spectra with and without the new molecular data. The results of these calculations can then be used to assess the importance of particular molecular opacities on the structure of the atmosphere. The synthetic spectra can be used to verify the correct strength of the computed bands when compared to observational data. This procedure introduces uncertainties such as the treatment of the equation of state (e.g., the molecular data used in it), the treatment of lines and line profiles and the assumed parameters of the comparison star. However, differential analyses circumvent many of these problems and should allow a reasonable evaluation of the molecular bound-free and free-free data.  The electronic bands of magnesium hydride have been detected over a wide range of stellar atmospheres including the photosphere of the sun \\citep{sot72}, sunspot umbrae \\citep{wal99}, F-K giants in the Milky Way halo and halos of other Local Group galaxies \\citep{maj00}, and nearby L-dwarfs \\citep{rei00}. MgH lines can be used as indicators of surface gravity in late-type stars \\citep{bon93} and to determine magnesium isotope abundances \\citep{wal99,gay00}.  The spectrum of MgH has been extensively studied in the laboratory for many decades \\citep[and references therein]{bal76,bal78,ber85,wal99} and has received some theoretical attention \\citep[and references therein]{sax78,kir79}. However, modern stellar atmosphere calculations require extensive, and complete, molecular line lists as molecular band absorption is the primary source of line-blanketing in cool stellar atmospheres, particularly M dwarfs. For many molecules, including MgH, the only source for complete line lists is the extensive compilations of \\citet{kur93}. While these compilations are highly valuable to stellar modelers, the methods necessary to compute 100s of millions of lines require a number of approximations which at times are severe. In this work, we apply fully quantum-mechanical techniques to compute the complete line list for the $A~^2\\Pi \\leftarrow X~^2\\Sigma^+$ transition of MgH. The parameters of the calculation are adjusted to force agreement with available experiments. However, our goal is to reproduce the global MgH opacity, hence we cannot claim spectroscopic accuracy for a particular line. The $A\\leftarrow X$ line list constructed in this work is combined with line lists for the $B'~^2\\Sigma^+ \\leftarrow X~^2\\Sigma^+$ and the $X~^2\\Sigma^+ \\leftarrow X~^2\\Sigma^+$ transitions computed by \\citet{sko02} and tested in a range of stellar atmosphere models. An overview of the theory of molecular rotational lines is presented in section 2 with the results of the line list calculations and stellar models are given in section 3. We present our conclusions in section 4. ",
        "conclusions": "Using a combination of theoretical and experimental data on the potential energies and dipole transition moment of MgH, a comprehensive theoretical vibrational-rotational line list for the $A~^2\\Pi \\leftarrow X~^2\\Sigma^+$ transition was constructed. When using the new $A\\leftarrow X$ line data and the new $B'~^2\\Sigma^+ \\leftarrow X~^2\\Sigma^+$ and the $X~^2\\Sigma^+ \\leftarrow X~^2\\Sigma^+$ line data of \\cite{sko02} in synthetic spectrum calculations, we find significant differences in the opacity when comparing the spectra to calculations using the existing data of \\citet{kur93}. The differences are largest for effective temperatures pertaining to L and M type stars and can easily be seen in low resolution work. For hotter stars, of K and G type, the differences are less pronounced and high resolution spectra are required to notice the improvements for the hottest stars."
    },
    "0206/astro-ph0206169_arXiv.txt": {
        "abstract": "In this contribution, I present the 90\\% confidence limits on the diagnostic diagrams of EW(WR bump) and L(WR bump)/L(H$\\beta$) ratio vs. EW(H$\\beta$) resulting from evolutionary synthesis models that include the statistical dispersion due to finite stellar populations in real star forming regions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206443_arXiv.txt": {
        "abstract": "We report on a deep search for radio pulsations toward five unidentified {\\it ASCA} X-ray sources coincident with {\\it EGRET} $\\gamma$-ray sources. This search has led to the discovery of a young and energetic pulsar using data obtained with the new Wideband Arecibo Pulsar Processor. PSR~J2021+3651 is likely associated with the X-ray source AX J2021.1+3651, which in turn is likely associated with the {\\it COS~B} high energy $\\gamma$-ray source 2CG~075+00, also known as GeV~J2020+3658 or 3EG~J2021+3716. PSR~J2021+3651 has a rotation period $P\\cong 104$~ms and $\\dot P \\cong 9.6\\times10^{-14}$, implying a characteristic age $\\tau_c\\sim  17$ kyr and a spin-down luminosity $\\dot E\\sim3.4\\times 10^{36}$ ergs~s$^{-1}$. The dispersion measure DM$\\simeq 371$~pc cm$^{-3}$ is by far the highest of any observed pulsar in the Galactic longitude range $55^{\\circ} < l < 80^{\\circ}$.  This DM suggests a distance $d\\ga 10$~kpc, and a high $\\gamma$-ray efficiency of $\\sim$ 15\\%, but the true distance may be closer if there is a significant contribution to the DM from excess gas in the Cygnus region. The implied luminosity of the associated X-ray source suggests the X-ray emission is dominated by a pulsar wind nebula unresolved by {\\it ASCA}. ",
        "introduction": "The majority of high energy $\\gamma$-ray sources observed by {\\it EGRET} and other telescopes have long escaped identification with lower energy counterparts \\citep{hbb+99}.  Young pulsars remain the only Galactic source class (other than the Sun) unambiguously shown to emit radiation in the 100~MeV -- 10~GeV range \\citep{tho01}. It is likely that many of the unidentified $\\gamma$-ray sources at low Galactic latitudes are young pulsars as well. Many of these sources have characteristics similar to those of the known $\\gamma$-ray pulsars, but have no known pulsars within their error boxes. This fact, along with modelling of the multi-wavelength pulse profiles and the still singular example of Geminga \\citep{hh92,hel94}, has led to the suggestion that a large fraction of the radio beams from $\\gamma$-ray sources will miss the Earth and appear radio quiet \\citep{rom96a}.  Recently, a number of young pulsars coincident with known $\\gamma$-ray sources have been discovered \\citep{dkm+01_mal,cbm+01}. These new discoveries are largely a result of greater sensitivity to pulsars with high dispersion measures (DM) obtainable with newer pulsar backends such as the Parkes multibeam system \\citep{mlc+01}. The recent detection of a young radio pulsar in the supernova remnant 3C58 with a 1400~MHz flux density of only $\\sim 50$~$\\mu$Jy \\citep{csl+02} suggests many more faint radio pulsars await discovery in deep, targeted searches.  A major stumbling block in the identification of the {\\it EGRET} sources is their large positional uncertainty, which can be greater than $1^{\\circ}$ across.  We approach this problem by targeting potential hard X-ray counterparts, whose size and positional uncertainty are much smaller than the typical single dish radio beam. Using as our guide the {\\it ASCA} catalog of potential X-ray counterparts of GeV sources (based on the \\citet{lm97} catalog of sources with significant flux above 1 GeV) by \\citet{rrk01} (hereafter, RRK), we have searched five X-ray sources for radio pulsations using the 305-m Arecibo telescope and the 64-m Parkes telescope (see Table~\\ref{tab:observations}). Previous searches of these targets were limited.  In particular, two of the three sources observed at Parkes (AX J1418.7$-$6058 and AX J1809.8$-$2332) were not previously the subject of any directed search and were only observed as a matter of course during the Parkes Multibeam Galactic Plane Survey \\citep{mlc+01}. A survey of {\\it EGRET} sources by \\citet{ns97} looked at two of the sources searched here (AX J1826.1$-$1300 and AX J2021.1+3651) with a limiting flux density for slow pulsars of 0.5 -- 1.0~mJy at frequencies of 370 and 1390~MHz, but found no new pulsars. Our search has led to the discovery of one young and energetic pulsar, PSR J2021+3658. We argue that it is a likely counterpart to AX J2021.1+3651 and GeV J2020+3658 / 2CG 075+00. ",
        "conclusions": "\\subsection{PSR J2021+3651}  Our search targeted AX J2021.1+3651, which was identified as a potential high-energy counterpart to GeV J2020+3658 by RRK. The X-ray source is near the {\\it ASCA} field edge and so the positional uncertainty from {\\it ASCA} is $\\ga 1^{\\prime}$ \\citep{guf+00}. A subsequent search of the {\\it ROSAT} All-Sky Survey Faint Source catalog \\citep{vab+00} revealed the source 1RXS J202104.5+365127 with a smaller positional error of $24^{\\prime \\prime}$. Given the rarity of such young, energetic pulsars and the small size of the Arecibo beam (3$^\\prime$ at FWHM), an association with the X-ray source is highly probable.  The DM of PSR J2021+3651 is by far the highest known in the Galactic longitude range $55^{\\circ} < l < 80^{\\circ}$ which is mainly an inter-spiral arm direction. The \\citet{tc93} model gives a distance of $\\sim 19$ kpc, well beyond the last spiral arm used in the model.  A revised model is currently in preparation by Cordes and Lazio which includes an outer spiral arm at $d\\sim 10$~kpc.  Placing the pulsar at the far edge of this outer arm still does not account for all the observed dispersion; and it is possible that there are further contributions from clouds in the Cygnus region, where there is known to be excess gas at $d\\sim 1.5$~kpc (J. Cordes, private communication).  However, there are no obvious HII regions within the Arecibo beam seen in either Very Large Array (VLA) 20-cm radio or Midcourse Space Experiment (MSX) 8.3~$\\mu$m images (available from the NASA/IPAC Infrared Science Archive).  The high DM is somewhat surprising given the X-ray absorption quoted by RRK, n$_{\\rm H}$=(5.0$\\pm0.25)\\times 10^{21}$~cm$^{-2}$, where the errors represent the 90\\% confidence region. The total Galactic HI column density in this direction as estimated from the FTOOL {\\it nh}, which uses the HI map of \\citet{dl90}, is $1.2\\times 10^{22}$~cm$^{-2}$.  This should be a good approximation if the source is truly at the far edge of the outer spiral arm. Noting that the {\\it ASCA} image shows faint, softer emission in the region (Figure~\\ref{fig:xray}), and given the likely possibility of either associated thermal X-ray flux from a supernova remnant or a nearby massive star, we fit the {\\it ASCA} spectrum of RRK, adding a thermal component to the absorbed power-law model. Accounting for $\\sim4$\\% of the photon flux with a MEKAL thermal plasma model of temperature $kT\\sim 0.1$~keV in XSPEC \\citep{arn96} statistically improves the fit ({\\it F}-test chance probability of 2.5\\%). The best-fit absorption for this three component model is n$_{\\rm H}$=7.6$\\times 10^{21}$~cm$^{-2}$ with a 90\\% confidence region of (4.1 -- 12.3)$\\times 10^{21}$~cm$^{-2}$, consistent with the total Galactic column density. The best-fit photon index is $\\Gamma=1.86$, still consistent with the 1.47 -- 2.01 range in RRK derived from the simple absorbed power-law model.  Hence the X-ray absorption does not force us to adopt a smaller distance than is suggested by the DM.  For a distance $d_{10}=d/10$~kpc, the inferred isotropic X-ray luminosity $L_X=4.8\\times 10^{34} d_{10}^2$ (2 -- 10~keV). The X-ray efficiency $\\eta_X=L_X/\\dot E$ is 0.01$d^2_{10}$. Compared to the total pulsar plus nebula X-ray luminosity of other spin-powered pulsars this is somewhat high, but within the observed scatter \\citep{pccm02,che00}.  The pulsar's positional coincidence with the error box of the hard spectrum, low variability {\\it EGRET} $\\gamma$-ray source GeV J2020+3658 coupled with the high inferred spin-down luminosity strongly suggests this pulsar emits pulsed $\\gamma$-rays. Unfortunately, confirming this by folding archival {\\it EGRET} data is problematic due to the likelihood of significant past timing noise and glitches, which make the back-extrapolation of the rotational ephemeris uncertain. RRK noted that the chance probability of an X-ray source as bright as AX J2021.1+3651 in the {\\it EGRET} error box was $\\sim10$\\%, but the nearby Wolf-Rayet star WR141 was equally bright in X-rays and also a potential $\\gamma$-ray emitter. However, young pulsars remain the only firmly established class of Galactic {\\it EGRET} sources. The known $\\gamma$-ray pulsars cluster at the top of pulsar lists rank-ordered by spin-down flux $\\dot E/d^2$, with $\\gamma$-ray efficiencies $\\eta_\\gamma=L_{\\gamma}/\\dot E$ mostly between 0.001 and 0.03 (assuming 1 sr beaming) with a tendency to increase with pulsar age \\citep{tbb+99}. The exception is PSR B1055$-$52, with an apparent $\\gamma$-ray efficiency $\\eta_\\gamma \\sim 0.2$ given its nominal DM distance of 1.5 kpc. The inferred $\\gamma$-ray efficiency for PSR J2021+3651 is $\\eta_\\gamma =0.15 d_{10}^2$ in the 100~MeV to 10~GeV range. If the pulsar is located within the Perseus arm at a distance of 5~kpc, then the inferred X-ray and $\\gamma$-ray luminosities would be fairly typical of the other pulsars with Vela-like spin-down luminosities. While there is currently no observational evidence for a distance this close, increased DM from an intervening source in this relatively crowded direction would not be surprising. We note that the DM derived distance for another young pulsar recently discovered within an $EGRET$ error box, PSR J2229+6114, also leads to an anomalously high inferred $\\gamma$-ray efficiency \\citep{hcg+01}.  \\subsection{Upper Limits Toward the Other Sources}  Determining the fraction of radio-quiet versus radio-loud pulsars is important for our understanding of $\\gamma$-ray pulsar emission mechanisms. The two leading classes of emission models, the outer-gap \\citep{rom96a} and polar-cap \\citep{dh96} models, make very different estimates of the fraction of $\\gamma$-ray pulsars that should be seen at radio energies. Out of the 25 brightest sources above 1 GeV not associated with blazars, $\\sim 10$ are now known to either be energetic radio pulsars or contain such pulsars within their error boxes. Searching the brightest unidentified X-ray sources in five GeV error boxes, we detected radio pulsations at the $\\sim 0.1$~mJy level (similar to the limiting sensitivity of the Parkes observations) from one of these with Arecibo. This is well below the average flux level expected for typical radio luminosities of young pulsars \\citep{bj99} and distances to star forming regions statistically associated with $\\gamma$-ray sources \\citep{yr97}. Two of the sources observed with Parkes, AX J1418.7$-$6058 (the Rabbit) and AX J1809.8$-$2333, have radio and X-ray properties that clearly identify them as pulsar wind nebulae \\citep{rrjg99,brrk02}, and the third, AX J1826.1$-$1300, is an extended hard X-ray source that has few other source class options. Therefore, all three remain viable candidates for $\\gamma$-ray loud, radio-quiet pulsars. Out of this same sample of 25 bright GeV sources, the total number of reasonable candidate neutron stars within the $\\gamma$-ray error boxes which have now been searched deeply for radio-pulsations without success is $\\sim 7$. A current ``best guess\" fraction of radio-loud $\\gamma$-ray pulsars of $\\sim 1/2$ falls in between the predictions of the two main competing models."
    },
    "0206/astro-ph0206396_arXiv.txt": {
        "abstract": "We have used the Arecibo telescope to carry out an survey of 31 dark clouds in the Taurus/Perseus region for narrow absorption features in HI ($\\lambda$ 21cm) and OH (1667 and 1665 MHz) emission. We detected HI narrow self--absorption (HINSA) in 77$\\%$ of the clouds that we observed. HINSA and OH emission, observed simultaneously are remarkably well correlated. Spectrally, they have the same nonthermal line width and the same line centroid velocity. Spatially, they both peak at the optically--selected central position of each cloud, and both fall off toward the cloud edges. Sources with clear HINSA feature have also been observed in transitions of CO, \\13co, \\c18o, and CI. HINSA exhibits better correlation with molecular tracers than with CI.  The line width of the absorption feature, together with analyses of the relevant radiative transfer provide upper limits to the kinetic temperature of the gas producing the HINSA. Some sources must have a temperature close to or lower than 10 K. The correlation of column densities and line widths of HINSA with those characteristics of molecular tracers suggest that a significant fraction of the atomic hydrogen is located in the cold, well--shielded portions of molecular clouds, and is mixed with the molecular gas.  The average number density ratio  [HI]/[\\h2] is $1.5\\times10^{-3}$. The inferred HI density appears consistent with but is slightly higher than the value expected in steady state equilibrium between formation of HI via cosmic ray destruction of H$_2$ and destruction via formation of H$_2$ on grain surfaces. The distribution and abundance of atomic hydrogen in molecular clouds is a critical test of dark cloud chemistry and structure, including the issues of grain surface reaction rates, PDRs, circulation, and turbulent diffusion. ",
        "introduction": "Two relatively distinct phases are generally assumed to exist in the neutral interstellar medium (ISM): atomic and molecular. The atomic phase of the ISM, consisting mainly of hydrogen atoms, is traced by the HI hyperfine transition at $\\lambda$ 21cm. The molecular phase of the ISM, whose major component -- molecular hydrogen -- lacks a permanent electric dipole moment and readily excited transitions at temperatures generally encountered, is primarily traced by emission from rarer molecular species such as carbon monoxide. The  conversion from atomic to molecular forms occurs on dust grains, where atomic hydrogen sticks and forms \\h2. This exothermic reaction releases \\h2 into the gas and keeps molecular clouds molecular. Dissociative processes maintain a population of atoms even inside molecular clouds. Recently, there has been increased interest in atomic species, which prove to be important probes. Two important examples are CI, which is accessible at submillimeter wavelengths \\citep*[e.g.\\ ][]{phillips80,huang99,plume00}, and OI, which can be observed in the far infrared \\citep*[e.g.\\ ][]{herr97, kram98, lise99, lis01}.  Inside molecular clouds, dissociating UV photons are blocked both by grains and by \\h2 line absorption \\citep{hws71}. A significant HI population exists inside molecular clouds maintained by cosmic ray destruction of \\h2 and additionally as a remnant of the \\h2 formation process in a chemically young cloud. The atomic hydrogen component inside molecular clouds has fractional abundance ([H]/[\\h2]) of $\\simeq$ 0.1$\\%$ (discussed in Section \\ref{nhih2}) and is thus the third most abundant gas phase species, after \\h2 and He. Because the balance of HI and \\h2 involves grain surface reactions, the density of atomic HI in dark clouds serves as a test of complete chemical networks with reactions both in the gas phase and on grain surfaces. The HI abundance in well--shielded regions can be increased by relatively rapid turbulent diffusion \\citep{willacy02}, or by general mass circulation \\citep{chie89}. It is therefore important to establish the presence of HI in the molecular ISM and to determine accurately its abundance.  A unique probe of this component is HI narrow (which we define to be less than that of CO) line width self--absorption, which we denote HINSA. The absorption dips seen in the spectra of $\\lambda$ 21cm emission are often denoted HI self--absorption. The prefix `self' is widely used to differentiate this phenomenon from absorption against a background continuum source. For the origins of HI absorption toward nearby dark clouds, \\emph{separate} galactic background HI emission and cold foreground HI material are both needed. A typical configuration is  shown in Figure~\\ref{fig:abs}.  The existence of cold HI associated with dark clouds was recognized more than 25 years ago. \\citet{knapp74} conducted a survey of 88 dark clouds and detected absorption features in fewer than half of them. The optical depth of cold HI was derived from the profiles of the absorption lines. The molecular content of the clouds observed was traced by their dust extinction. The HI fractional abundance, [H]/[\\h2] was thus determined to be 1\\% to 5\\%. In terms of the observational limits and conclusions, this early work typifies most self--absorption studies that have followed.  With the 140ft (43m) telescope, Knapp's survey has an angular resolution of 21\\arcmin\\ and a velocity resolution of $\\sim$0.5 \\kms. Similar resolutions have been achieved with the 85ft (26m) antenna at Hat Creek \\citep{goodman94} and the 120ft (37 m) antenna of the Haystack Observatory (Myers et al. 1978). The 76m Lovell telescope (12\\arcmin\\ angular resolution, $\\sim$0.5 \\kms\\ velocity resolution; the same convention in what follows) has been employed for a study of six dark clouds in the Lynds (1962) catalogue (McCutcheon, Shuter \\& Booth  1978) and the Riegel--Crutcher cloud (Montgomery, Bates \\& Davies 1995). Additional studies have been conducted with the Effelsberg 100m telescope (9\\arcmin, 0.5 \\kms). The Taurus molecular cloud TMC1 has been mapped by Wilson and Minn (1977). The complex region around B18, also known as Kutner's cloud, has been mapped by Batrla, Wilson \\& Rache (1981) and by P\\\"{o}ppel, Rohlfs \\& Celnik (1983). The Arecibo telescope with the line feed system (3\\arcmin, 1 \\kms) has been used to study HI absorption  \\citep*[e.g.\\ ][]{burton78, baker79, bania84}. Baker \\& Burton (1979) also raise the possibility that the absorption can help resolve the near--far ambiguity in kinematic distances, a topic which is further explored by Jackson et al.\\ (2002).  Interferometers have also been used to obtain higher angular resolution, but the usual penalty has been lower velocity resolution to augment the sensitivity. Van der Werf et al.\\ mapped L134 (1988) and L1551 (1989) with the DRAO and the VLA (1.5\\arcmin, 1.3 \\kms). The interferometer studies are well--suited for mapping structures having a scale of arc--minutes. But their velocity resolutions of about 1.5 \\kms\\ can easily miss or suppress narrow absorption features, which occupy at most a couple of velocity channels in such spectra. To compound this problem, the HI emission is structured. Multiple peaks and wide troughs are not rare in galactic 21c profiles (see, e.g. Figure~\\ref{dipoff}). If a map is based on the integrated area of an absorption line (or lines), the wide troughs which may not be associated with dark clouds will be more prominent than are the narrow features. Reliable analysis of HINSA profile requires a velocity resolution better than 0.3 \\kms.  The general scientific objectives of narrow line HI absorption studies are to explain its origin and to determine the abundance of atomic hydrogen. For the first question, the association of HINSA with dark clouds is inconclusive in the literature. On the positive side, Sherwood \\& Wilson (1981) find a good correlation with extinction in TMC1. McCutcheon et al.\\ (1978) give a higher detection rate (6--8 out of 11 Lynds clouds) than Knapp (1974). Cappa de Nicolau \\& Poppel (1991) find HINSA in the `darkest' cores in the CrA complex embedded in a HI emission ridge. On the other hand, the association between HI  absorption features and dense regions is ambiguous in a recent DRAO HI survey, in which Gibson et al.\\ (2000) find cloud--like absorption structures both correlated and uncorrelated with CO emission. They label these features HISA, shorthand for HI self--absorption. This is the situation in which the distinction between HISA and HINSA must be made. As seen in the DRAO survey, the HISA  probably reflects temperature fluctuations in the atomic ISM. The HISA is spectrally wider and flatter and could have its origin in the same location as the 21cm emitting gas, making its name appropriate. HINSA refers to narrower absorption features, which are produced by cold foreground molecular clouds and which are not prominent in the DRAO survey with 1.3 \\kms\\ velocity resolution. To further complicate the issue, HI has also been seen in {\\em emission} in possible halos around molecular clouds, such as B5 \\citep{wannier91}. HI halos are a distinct environment for atomic hydrogen compared to either HISA or HINSA, and they are easily distinguished observationally (see Section~\\ref{discussion}).  The difficulties in obtaining the HI abundance through HINSA are three--fold. First, the early studies are hampered mainly by limited knowledge of the dark cloud itself. The  \\h2 column density is often obtained from low angular resolution extinction data (e.g.\\ Knapp 1974, Batrla et al.\\ 1981) or H$_2$CO (e.g. Wilson \\& Minn 1977, Poppel et al.\\ 1983), a molecule whose fractional abundance is not very certain. Second, with one profile, it is impossible to obtain both the optical depth and the spin temperature accurately. An assumption about the spin temperature is often made, which may not be realistic. Third, the issue of foreground emission is often ignored.  With some or all of the three uncertainties, the [H]/[\\h2] ratio has been derived to be ranging from a few percent (e.g.\\ Knapp 1974; Saito, Ohtani \\& Tomita 1978), to 5$\\times$10$^{-4}$ \\citep{winnberg80}.  The requirements of reasonable angular resolution, good sensitivity, and high frequency resolution make the upgraded Arecibo Gregorian system a valuable tool with which to study the HINSA. The large instantaneous bandwidth allows OH 1665 and 1667 MHz spectra to be obtained simultaneously with that of HI. We describe the new observations including a HI survey to examine the correlation between HINSA and OH in dark clouds and complementary mapping in carbon monoxide isotopologues, OH, and CI in Section 2. In Section 3 we present a three--component model for radiative transfer and correction for foreground material, allowing accurate HI column densities to be obtained. We analyze OH and \\c18o emission in Section 4, and the results of our survey and observations of L1544 in Section 5. We review the general issue of atomic hydrogen in molecular clouds in Section 6, discuss our results in Section 7, and summarize our conclusions in Section 8. ",
        "conclusions": "We have surveyed 31 dark clouds using Arecibo, FCRAO, and SWAS. The analysis of these data show  the following. \\begin{enumerate}  \\item The  21cm HI narrow self--absorption (having line width smaller than that of the corresponding CO emission) is a widespread phenomenon, detected in $\\simeq$ 77$\\%$ of our sample of dark clouds in the Taurus/Perseus region. We use the term HINSA to distinguish the narrow absorption definitely caused by molecular cooling from broader absorption features seen in other surveys of HI throughout the Galaxy.  \\item The atomic hydrogen producing the HINSA absorption has significant column density with $N(\\text{HINSA}) \\sim 7\\times 10^{18}$ \\cm2.  \\item The gas responsible for the HINSA is at low temperatures, between 10 and 25 K. Some sources (L1521E, L1512, L1523) must be thermalized at temperatures close to or lower than 10 K.  \\item The nonthermal line width of HINSA is comparable to the line width of \\13co, only slightly larger than that of \\c18o, and is smaller than the line widths of CO or CI. This suggests that HINSA is produced by cold atomic hydrogen in regions of moderate extinctions with $A_v$ larger than a few.  \\item In the maps of L1544, HINSA is morphologically similar to \\c18o.  \\item The low temperature, the absence of increased absorption at cloud edges, and the narrow line width of HINSA suggest that the atomic hydrogen producing HINSA is mixed with the gas in cold, well--shielded regions of molecular clouds. \\end{enumerate}"
    },
    "0206/astro-ph0206449_arXiv.txt": {
        "abstract": "We report the detection of fully resolved absorption lines of {\\it A$-$X} bands from interstellar $^{12}$C$^{17}$O and $^{12}$C$^{18}$O, through high-resolution spectroscopy of \\objectname{X Per} with the Space Telescope Imaging Spectrograph\\footnotemark[3]. The first ultraviolet measurement of an interstellar $^{12}$C$^{17}$O column density shows that its isotopomeric ratio is $^{12}$C$^{16}$O/$^{12}$C$^{17}$O = 8700 $\\pm$ 3600. Simultaneously, the second ultraviolet detection of interstellar $^{12}$C$^{18}$O establishes its isotopomeric ratio at 3000 $\\pm$ 600. These ratios are about five times higher than local ambient oxygen isotopic ratios in the ISM. Such severe fractionation of rare species shows that both $^{12}$C$^{17}$O and $^{12}$C$^{18}$O are destroyed by photodissociation, whereas $^{12}$C$^{16}$O avoids destruction through self-shielding. This is to be contrasted with our ratio of $^{12}$C$^{16}$O/$^{13}$C$^{16}$O = 73 $\\pm$ 12 toward X Per, which is indistinguishable from $^{12}$C/$^{13}$C, the result of a balance between photodissociation of $^{13}$C$^{16}$O and its preferential formation via the isotope exchange reaction between CO and C$^+$. ",
        "introduction": "\\footnotetext[3] {Based on observations obtained with the NASA/ESA {\\it Hubble Space Telescope (HST)} through the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555.}  Carbon monoxide is the second most abundant molecule in interstellar clouds after H$_2$, with readily observable electronic transitions in the vacuum ultraviolet (VUV), vibrational bands in the infrared, and pure rotational lines in the mm-wave regime. The isotopic varieties of CO are used to constrain models of star formation, chemical networks, and stellar evolution. The first VUV detection of rotationally unresolved fourth-positive ($A~^1\\Pi$$-$$X~^1\\Sigma^+$) absorption bands from interstellar $^{12}$C$^{16}$O and $^{13}$C$^{16}$O was reported by \\citet{ss}--see \\citet{mn} for an extensive review of the fourth-positive bands. However, most measurements of various CO species come from radio observations of molecular clouds with substantial total column densities ($N$) of CO. \\citet{p72} reported the first mm observations of $^{12}$C$^{16}$O and $^{13}$C$^{16}$O in dark clouds. Rarer CO varieties were reported later: $^{12}$C$^{18}$O by \\citet{mms}, $^{12}$C$^{17}$O by \\citet{d77}, and $^{13}$C$^{18}$O by \\citet{l80}. The ``final'' milestone for radio CO was recently reached by \\citet{b01}, who reported the detection of the rarest stable CO isotopomer, $^{13}$C$^{17}$O, in the $\\rho$ Oph molecular cloud. In the VUV, absorption from CO is sought, but such observations sample substantially smaller $N$(CO) than are commonly observed in the radio because ultraviolet extinction limits the number of suitable targets behind molecular clouds. Individual {\\it A$-$X} rotational lines of interstellar absorption bands have been fully resolved and measured only for $^{12}$C$^{16}$O and $^{13}$C$^{16}$O, using the echelle grating of the Goddard High-Resolution Spectrograph (GHRS) on board the $HST$ \\citep{s91,s92}. A previous detection with a GHRS first-order grating of rotationally unresolved $^{12}$C$^{18}$O was reported by \\citet{l94} toward $\\zeta$ Oph.  In this Letter we report the first VUV detection of interstellar absorption from $^{12}$C$^{17}$O bands, and present the first unblended measurements of its sibling, $^{12}$C$^{18}$O. These observations were made toward X Per (HD 24534), using grating E140H of the Space Telescope Imaging Spectrograph (STIS) for data sets o64812010$-$030 and o64813010$-$020. The star was observed through the smallest aperture (0.1\\arcsec~$\\times$ 0.025\\arcsec, the ``Jenkins slit''), providing the highest $HST$ resolving power of $\\lambda/\\Delta\\lambda$ = 200,000 over $\\lambda$ = 1316 to 1517 \\AA. We modeled the data using our unpublished spectrum synthesis code, ISMOD, which is based on the line transfer equations given by \\citet{bvd}. For more details of data extraction and modeling, see \\citet{sfl}. In the next section we present CO column densities along the X Per line of sight, and in \\S 3 we compare CO isotopomeric ratios with ambient carbon and oxygen isotopic ratios in the context of theoretical models of translucent clouds. ",
        "conclusions": "The line of sight toward X Per has provided us with a rich spectrum of {\\it A$-$X} bands of CO, allowing the first VUV detection of the rare isotopomer $^{12}$C$^{17}$O and the first rotationally-resolved views of both $^{12}$C$^{18}$O and $^{12}$C$^{17}$O. These detections were made possible by the superb qualities of STIS as a VUV spectrometer. Toward X Per we find that $^{13}$C$^{16}$O is unfractionated with respect to $^{12}$C$^{16}$O owing to a balance between the rates of photodissociation and of the isotope exchange reaction. On the other hand, the lack of an isotope exchange reaction in the case of oxygen isotopes renders both $^{12}$C$^{18}$O and $^{12}$C$^{17}$O strongly fractionated and destroyed at the 80\\% level with respect to the strongly shielded $^{12}$C$^{16}$O. As was described above, similar isotopomeric ratios are to be found within published results from theoretical models. However, a large gap remains between observed and modeled column densities for CO.  With the detection of $^{12}$C$^{17}$O toward X Per, observers of interstellar absorption lines are left with the last two stable CO isotopomers that have yet to be detected in the ultraviolet, namely, $^{13}$C$^{18}$O and $^{13}$C$^{17}$O. These are very challenging tasks at best, as the respective intrinsic abundances are expected to be \\case{1}{24} and \\case{1}{70} the abundance of $^{12}$C$^{17}$O. Whereas $^{13}$C$^{18}$O was included in models of translucent clouds, not even a coarse grid of theoretical models exists for $^{12}$C$^{17}$O in translucent clouds. Hopefully, the new observations of this species as reported here and of $^{13}$C$^{17}$O by Bensch et al. (2001) will provide an incentive for the inclusion of $^{17}$O-bearing molecules in computer simulations. There is also a need for laboratory measurements, since the {\\it A$-$X} bands of $^{12}$C$^{17}$O currently lack rigorous wavelength and perturbation analyses."
    },
    "0206/astro-ph0206163_arXiv.txt": {
        "abstract": "In an ongoing infrared imaging survey of quasars at Keck Observatory, we have discovered that the $z=1.285$ quasar SDSS~J233646.2-010732.6 comprises two point sources with a separation of 1\\farcs67.  Resolved spectra show that one component is a standard quasar with a blue continuum and broad emission lines; the other is a broad absorption line (BAL) quasar, specifically, a BAL QSO with prominent absorption from \\MgII\\ and metastable \\FeII, making it a member of the ``FeLoBAL'' class.  The number of known FeLoBALs has recently grown dramatically from a single example to more than a dozen, including a gravitationally lensed example and the binary member presented here, suggesting that this formerly rare object may be fairly common. Additionally, the presence of this BAL quasar in a relatively small separation binary adds to the growing evidence that the BAL phenomenon is not due to viewing a normal quasar at a specific orientation, but rather that it is an evolutionary phase in the life of many, if not all, quasars, and is particularly associated with conditions found in interacting systems. ",
        "introduction": " ",
        "conclusions": "A deep imaging survey of BAL quasars and a control sample of non-BAL quasars, from the ground or space, is needed to document the frequency of each in systems that show tidal interactions, which would be a more comprehensive version of the work done by Canalizo \\& Stockton (2001). If BAL-ness is explained simply by viewing an ordinary quasar along a line of sight which skims the ``dusty torus'' (Weymann et al.\\ 1991), then the two populations should appear with equal frequency in merging or interacting systems.  If BAL features are produced by conditions created in mergers, and orientation plays less or no role, then they will be more common in chaotic systems with tidal tails and other signs of host galaxy interactions.  Dynamical modeling of the tidal effects may help constrain the timescales of various quasar phases.  FeLoBALs may be a common feature of the AGN landscape, especially if there are Seyfert-type luminosity class analogs to the brighter objects now being unearthed.  If FeLoBALs are indeed common, then infrared sky surveys which reach deeper than 2MASS, such as those being undertaken with SIRTF, will turn up numerous examples along with other extreme BAL phenomena in quasar and galaxy luminosity objects. The FeLoBAL-like Hawaii~167 (Cowie et al.\\ 1994) was discovered in just such a survey.  Such surveys are needed to provide more accurate estimates of the extremely red quasar population, especially at redshifts above 2.5 where such objects are essentially invisible in the optical."
    },
    "0206/astro-ph0206480_arXiv.txt": {
        "abstract": "{We report on an \\xmm\\ observation of the X-ray afterglow of the Gamma Ray Burst GRB 011211, originally detected by \\sax\\ on 11th December 2001. The early afterglow spectrum obtained by \\xmm, observed 11 hours after the initial burst, appeared to reveal decaying H-like K$\\alpha$ emission lines of Mg, Si, S, Ar and Ca, arising in enriched material with an outflow velocity of order 0.1c (Reeves \\et 2002). This was attributed to matter ejected from a massive stellar progenitor occurring shortly before the burst itself. Here, we present a detailed re-analysis of the \\xmm\\ EPIC observations of GRB 011211. In particular, we show that the detection of the soft X-ray line emission appears robust, regardless of detector background, calibration, spectral binning, or the spectral model that is assumed. We demonstrate that thermal emission, from an optically thin plasma, is the most plausible model that can account for the soft X-ray emission, which appears to be the case for at least two burst afterglow spectra observed by \\xmm. The X-ray spectrum of GRB 011211 evolves with time over the first 12 ksec of the \\xmm\\ observation, the observations suggest that thermal emission dominates the early afterglow spectrum, whilst a power-law component dominates the latter stages. Finally we estimate the mass of the ejected material in GRB 011211 to be of the order 4-20 solar masses.} ",
        "introduction": "The study of X-ray afterglow emission is crucial to understanding the nature of gamma ray bursts and their progenitors, as X-ray spectroscopy can reveal details of the environment of the burst explosion. For instance, observations of some X-ray afterglows with \\sax, \\asca, and {\\it Chandra} have revealed strong iron K$\\alpha$ emission lines (e.g. GRB 000214, Antonelli \\et 2000; GRB 991216, Piro \\et 2000; GRB 970828, Yoshida \\et 1999; GRB 970508, Piro \\et 1998), with line equivalent widths of up to several keV; typically the lines are detected at $\\sim3\\sigma$ confidence. An iron K-shell absorption feature was also reported in the prompt \\sax\\ spectrum of one burst, GRB 990705 (Amati \\et 2000, Lazzati \\et 2001). The line observations appear to support models where the burst explodes into an enriched, high density medium, favoring a massive stellar progenitor for long duration bursts. In many cases (e.g. Piro \\et 2000, Antonelli \\et 2000), large masses of iron are required to account for the line emission. One possible source of the iron is in a shell of distant (R~$\\sim10^{16}$~cm) material ejected from a supernova (the `supranova' model, Vietri \\& Stella 1998) which would have to occur at least several months {\\it prior} to the burst event, in order to allow sufficient time for large quantities of iron to form (e.g. Vietri \\& Stella 1998). Alternatively, the iron K emission could arise through the interaction of a magnetically driven wind with the envelope of the massive progenitor star (Rees \\& Meszaros 2000); here the distances involved are much smaller (R$\\sim10^{13}$~cm) and consequently the high masses of iron and long time delays between supernova and burst are not required.  The high throughput of \\xmm, compared with \\sax\\ or {\\it Chandra}, makes it the best available telescope with which to constrain X-ray line emission in GRBs, potentially providing a powerful diagnostic for discriminating between models of the burst progenitor. In an analysis of the \\xmm\\ observation of the afterglow of gamma ray burst GRB 011211, Reeves \\et (2002) reported the first detection of emission lines other than iron in a GRB afterglow spectrum; specifically the decaying line emission from the hydrogenic states of Mg, Si, S, Ar and Ca. Furthermore the energies of the X-ray lines appeared to be offset from the known redshift for the host galaxy of GRB 011211 ($z=2.140\\pm0.001$, Holland \\et 2002), implying the line emitting matter was outflowing with a velocity of $\\sim$~0.1c from the site of the GRB progenitor. This result was interpreted as possible evidence of a supernova explosion occurring within days of the burst itself, with the X-ray line emission arising from matter in the expanding supernova shell.  Recently, Borozdin \\& Trudolyubov (2002) have claimed that the detection of soft X-ray emission lines from light metals in GRB 011211 may not be robust, as the observations could be contaminated by the background of the \\xmm\\ EPIC-pn detector. In addition, Rutledge \\& Sako (2002) perform simulations which suggest that the significance of the emission features may be lower, at the level (at best) of $\\sim98.8\\%$ confidence. In this paper, we describe our analysis of the \\xmm\\ data on GRB 011211 in detail. We confirm that the soft X-ray emission features in GRB 011211 are robust, the set of lines being detected to a good level of confidence (at $>99$\\% significance in all our tests). Both the F-test and Monte-Carlo simulations are used to determine the statistical significance of the line emission features, making no prior assumption about the rest energies or redshift of the line emitting material. We also demonstrate that the results do not depend on the background subtraction of the EPIC-pn detector, instrument calibration, or the spectral binning used and show that the result is independent of the spectral model that is assumed. In section 2 we review the properties of the burst, the \\xmm\\ observations and basic data analysis and in section 3 detail the spectral analysis of the afterglow during the first 5 ksec when the lines appear most prominent. The time dependent spectral properties of GRB 011211 are also outlined, whilst a likely mechanism of the line emission is discussed and a mass for the ejected material in GRB 011211 is derived. Throughout this paper we adopt a cosmology of $H_{0}= 75$~km~s$^{-1}$~Mpc$^{-1}$ and $q_{0}=0.1$. Unless otherwise stated, all errors are quoted at 68\\% confidence (i.e. $\\Delta\\chi^2=1.0$ or 2.3 for 1 or 2 interesting parameters respectively). ",
        "conclusions": "Both the F-test and Monte-Carlo simulations have showed that the soft X-ray line emission in GRB 011211 is detected at $>99\\%$ confidence, We conclude that the detection of the soft X-ray features in GRB 01211 is robust, and is not affected by the detector background, calibration, spectral binning or the particular spectral model that is assumed. It is also apparent that the lines decay rapidly, no line emission is formally detected after 10~ks from the start of the \\xmm\\ observation.  The early afterglow data (during the first 10 ksec of observation) appears to be characterised by thermal X-ray emission (from optically thin gas), as evident by the curvature present in the in the early spectrum. An alternative spectral model, where the afterglow line emission results from X-ray reflection off optically thick material appears to be ruled out (e.g. Ballantyne \\& Ramirez-Ruiz 2001), although we note that in current reflection models, the emission from S, Ar and Ca is not computed. Future models will hopefully provide a better test of the reflection scenario.  On the assumption that the lines result from thermal emission one can estimate the mass of the ejected material. Firstly the emission measure ($E_{M}$) is computed from the observed luminosity of the thermal emission component. Here $n^{2}V=E_{M}\\sim10^{69}$cm$^{-3}$ (where n is the electron density in cm$^{-3}$ and V is the volume), whilst the plasma will emit X-rays over a characteristic cooling time (governed by the electron density) of $t_{cool}=1.4\\times10^{15}n^{-1}$~secs, for gas of temperature $kT=4$~keV. Thus from the emission measure, the total integrated output of the thermal emission is governed by $n^{2}V=E_{M}t_{line}/t_{cool}$, where $t_{line}$ is the observed lifetime of the line emission. Substituting for the cooling time, and as $nV=M/m_{p}$ (where M is the total mass of the emitting material and $m_{p}$ is the proton mass), then one can derive the total mass of the ejecta:-  \\begin{center}  $M = 2.4\\times 10^{30} t_{line}$~g \\\\ or $M = 1.2 \\times 10^{-3} t_{line}~M_{\\odot}$  \\end{center}  The minimum lifetime of the line emission (measured from the start of the \\xmm\\ observation) is $10^{4}/(1+z)$~secs, whilst the maximum lifetime occurs if the line emission starts at the time of the burst, which is then $t_{line}=5.5\\times10^{4}/(1+z)$~secs. Thus the minimum and maximum mass of the outflowing material, responsible for the blue-shifted line emission, are then 4 and 20 solar masses respectively. This indicates that the mass of the stellar progenitor is likely to be $>20M_{\\odot}$ even for the conservative case where the mass of the {\\it ejecta} is $4M_{\\odot}$ (e.g. Woosley \\& Weaver 1995).  The duration of the line features in GRB 011211, observed 11 hours after the initial burst, implies a distance of several~$\\times~10^{15}$~cm for the line emitting matter (Reeves \\et 2002). When combined with the velocity of the material ($\\sim0.1c$), this implies that the matter was ejected from days to weeks before the burst itself occurred. This time delay appears too short in the context of the ``supranova'' model (Vietri \\& Stella 1998), where a time-delay of several months to years can be predicted, whilst the ``hypernova'' class of models (Woosley 1993) involve the near-simultaneous occurrence of a supernova and the gamma ray burst. Interestingly one model has been proposed involving the coalescence of a close binary system, resulting from a collapse of a massive stellar progenitor (Davies \\et 2002). This could naturally account for the time delay between the putative supernova and the gamma ray burst, with a delay of a few days resulting from the orbital decay (through gravitational radiation) of the close binary system.  If the line emitting material is located at much smaller distances ($R\\sim10^{13}$~cm), as in nearby reprocessor models (e.g. Rees \\& Meszaros 2000) then the need for a significant time delay is reduced. Here, the line emission results from re-processing (via reflection/photoionisation) of continuum photons emitted {\\it after} the initial burst. However in this scenario, significant continuum emission should be observed at the same time as the line emission, which would render any soft X-ray line features undetectable by current instrumentation (see Lazzati \\et 2002). One possibility that has been suggested (Kumar \\& Narayan 2002) is that the ionising photons are scattered back towards the ejected material, by a positron-electron shield generated by the initial $\\gamma$-rays themselves. This model removes the need for a supernova-GRB delay, as the duration of the line emission is accounted for by the time required for the photons to be scattered back towards the ejected matter.  Including GRB 011211, there have now been 4 detections of GRB afterglows with \\xmm\\ (GRB 001025A and GRB 010220, Watson \\et 2002; GRB 020322, Ehle \\et 2002), whilst the X-ray afterglow for the faint burst GRB 020321 could not be localised. Interestingly, line emission features are also indicated in the X-ray afterglows of GRB 001025A and GRB 010220 (Watson \\et 2002). Similar to GRB 011211, the spectrum of GRB 001025A appears to exhibit a blend of soft X-ray line emission, from the medium-z elements such as Mg, Si, S, and Ar. In contrast, the spectrum of GRB 010220 shows a strong iron-group (i.e. Fe, Co or Ni) line (equivalent width $\\sim1$~keV), although here the X-ray spectrum is absorbed below 1 keV, rendering any soft X-ray line emission undetectable. Both fits favour an over abundance of Nickel (or Cobalt) to iron, which is suggestive of a short time delay (days rather than months) between the putative supernova and burst. Thermal emission models are formally required to fit the soft X-ray excess observed in GRB 001025A, whilst the spectrum of GRB 010220 can be equally well fitted by either thermal or reflection models. Thus soft X-ray line emission is required in at least 2 of the \\xmm\\ afterglow spectra (GRB 011211 and GRB 001025A), the exception being the recent burst GRB 020322, which  appears to exhibit a featureless (but absorbed) power-law spectrum (Reeves \\et 2002b, in preparation).  It seems possible that soft X-ray line emission features are relatively common in X-ray afterglow spectra, whilst thermal emission models can fit 3 out of the 4 present \\xmm\\ spectra. The previous lack of detection of soft X-ray features with Chandra and \\sax\\ can be accounted for by the lower effective area of the ACIS-S and LECS detectors respectively in the soft X-ray band. Indeed \\xmm\\ is the first X-ray mission with the sensitivity required to detect any putative soft X-ray features in GRBs at energies below 2 keV. It is important that \\xmm\\ performs follow-up observations of bright ($>10^{-13}$~erg~cm$^{-2}$~s$^{-1}$) afterglows, within several hours of the burst, to determine the frequency the soft X-ray line emission. Ultimately, through monitoring the temporal behavior of the line emission from minutes to many hours after a burst, the X-ray telescope on-board NASA's forthcoming {\\it Swift} mission will be able to trace the distribution of matter around the burst explosion and solve the ambiguities that arise when discriminating between the current GRB models."
    },
    "0206/astro-ph0206213_arXiv.txt": {
        "abstract": "{ { In this article we present the case of HD\\,41004\\,AB, a system composed of a K0V star and a 3.7-magnitude fainter M-dwarf companion. We have obtained 86 CORALIE spectra of this system with the goal of obtaining precise radial-velocity measurements. Since HD\\,41004\\,A and B are separated by only 0.5\\arcsec, in every spectrum taken for the radial-velocity measurement, we are observing the blended spectra of the two stars. An analysis of the measurements has revealed a velocity variation with an amplitude of about 50\\,m\\,s$^{-1}$ and a periodicity of 1.3\\,days. This radial-velocity signal is consistent with the expected variation induced by the presence of a companion to either HD\\,41004\\,A or HD\\,41004\\,B, or to some other effect due to e.g. activity related phenomena. In particular, such a small velocity amplitude could be the signature of the presence of a very low mass giant planetary companion to HD\\,41004\\,A, whose light dominates the spectra. The radial-velocity measurements were then complemented with a photometric campaign and with the analysis of the bisector of the CORALIE Cross-Correlation Function (CCF). While the former revealed no significant variations within the observational precision of $\\sim$0.003-0.004\\,mag (except for an observed flare event), the bisector analysis showed that the line profiles are varying in phase with the radial-velocity. This latter result, complemented with a series of simulations, has shown that we can explain the observations by considering that HD\\,41004\\,B has a brown-dwarf companion orbiting with the observed 1.3-day period. As the spectrum of the fainter HD\\,41004\\,B ``moves'' relative to the one of HD\\,41004\\,A (with an amplitude of a few km\\,s$^{-1}$), the relative position of the spectral lines of the two spectra changes, thus changing the blended line-profiles. This variation is large enough to explain the observed radial-velocity and bisector variations, and is compatible with the absence of any photometric signal. If confirmed, this detection represents the first discovery of a brown dwarf in a very short period (1.3-day) orbit around an M dwarf. Finally, this case should be taken as a serious warning about the importance of analyzing the bisector when looking for planets using radial-velocity techniques.} ",
        "introduction": "{ Radial-velocity techniques have so far unveiled about 80 planetary companions around} solar type dwarfs\\footnote{See e.g. obswww.unige.ch/$\\sim$udry/planet/planet.html}. The most precise instruments currently available for planet searches can measure the velocity of a star in the direction of the line-of-sight with a precision of the order of 2-3\\,m\\,s$^{-1}$ \\citep[e.g.][]{Que01a,But01,Pep02}, but even higher precision is expected from instruments available in the near future \\citep[e.g. HARPS -- ][]{Pep00}. { This will definitely allow the discovery of lower mass and longer period planets}, that remained undetected up to now due to the low amplitude of the induced radial-velocity variation.  The gain in precision will, however, bring to light some of the limitations of the radial-velocity method. It is well known, for example, that radial-velocity ``jitter'' with amplitudes up to a few tens of m\\,s$^{-1}$ is expected to result from the presence of strong photospheric features like spots or convective inhomogeneities, associated with chromospheric activity phenomena \\citep[][]{Saa97,Saa98,San00}. The presence of spots can even induce a periodic radial-velocity signal similar to the one expected from the presence of a planet. This is the case for \\object{HD\\,166435} \\citep[][]{Que01b}, a star presenting a radial-velocity signal with a period of about 3.8-days, but showing both photometric and bisector variations with the same periodicity.  \\begin{table} \\caption[]{ Stellar parameters for \\object{HD\\,41004\\,A}} \\begin{tabular}{lcc} \\hline \\hline \\noalign{\\smallskip} Parameter  & Value & Reference \\\\ \\hline \\\\ $Spectral~type$  & K1V/K2V & Hipparcos \\citep[][]{ESA97}/  \\\\ &         & {\\it uvby} (see Sect\\,\\ref{sec:photom})  \\\\ $Parallax$~[mas]  & 23.24 $\\pm$ 1.02 & Hipparcos \\citep[][]{ESA97}\\\\ $Distance$~[pc]  & 43 & Hipparcos \\citep[][]{ESA97} \\\\ $m_v$  & 8.65 & Hipparcos \\citep[][]{ESA97}\\\\ $B-V$  & 0.887 & Hipparcos \\citep[][]{ESA97} \\\\ $T_\\mathrm{eff}$~[K]  & 5010 & See text \\\\ $\\log{g}$~[cgs]  & 4.42 & See text \\\\ $M_\\mathrm{v}$  & 5.48 & -- \\\\ $Luminosity~[L_{\\sun}]$ & 0.65 & \\citet{Flo96} \\\\ $Mass~[M_{\\sun}]$  & $\\sim$0.7 & -- \\\\[5.0pt] $\\log{R'_\\mathrm{HK}}$  & $-$4.66 & \\citet[][]{Hen96} \\\\ $Age$~[Gyr] & 1.6 & \\citet[][]{Don93} \\\\ $P_{\\mathrm{rot}}$~[days]   & $\\sim$27 &  \\citet[][]{Noy84} \\\\[5.0pt] $v\\,\\sin{i}$~[km\\,s$^{-1}$] & 1.22 & CORALIE \\\\ $\\mathrm{[Fe/H]}$  & $-$0.09/$+$0.10 & {\\it uvby}/CORALIE  \\\\ \\hline \\noalign{\\smallskip} \\end{tabular} \\label{tab1} \\end{table}  The case of \\object{HD\\,166435} illustrates very well the need to confirm, at least for the shortest period cases, that the radial-velocity signature is indeed due to the presence of a low mass companion, and not due to some kind of intrinsic phenomena. As shown by \\citet[][]{Que01b}, the use of photometric data and bisector analysis was crucial to clarify the origin of the radial-velocity variations observed on this star.  In this paper we present the case of \\object{HD\\,41004}, a visual double system consisting of a K1V-M2V pair (A and B components). This system was found to present a radial-velocity signature similar to the one expected as if the K1 dwarf had a very low mass planetary companion in a 1.3-day period orbit. Although the photometric data revealed no significant photometric variations, an analysis of the Bisector Inverse Slope (BIS) of the CORALIE Cross-Correlation Function (CCF) \\citep[][]{Que01b} revealed a periodic variation in phase with the radial-velocity signal. In the following sections we will show that the radial-velocity variation is in fact not a result of the periodic motion of the A component, but of the Doppler motion of the spectrum of the B component due to the presence of a brown-dwarf companion. The results strongly caution about the need to use methods capable of detecting line asymmetries, like the bisector analysis, when dealing with high-precision planet searches with radial-velocity techniques. ",
        "conclusions": "{ We have presented the case of \\object{HD\\,41004\\,AB}, a system composed of a K0V star and a 3.7 magnitudes fainter M2 dwarf, separated by 0.5\\arcsec. Radial-velocity measurements derived from CORALIE blended spectra of the two stars have unveiled a radial-velocity variation with a period of $\\sim$1.3\\,days and a small amplitude ($\\sim$50\\,m\\,s$^{-1}$), compatible with the expected signal due to the presence of a planetary companion to \\object{HD\\,41004\\,A}. However, as we have seen, the combined radial-velocity, photometry and bisector analysis suggest that the best explanation for the observations is that the fainter \\object{HD\\,41004\\,B} has a brown-dwarf companion. In this scenario, the observed low amplitude radial-velocity variation is due to the measure of a variation in the line profiles which is induced by the change of the relative position of the spectra corresponding to the two stellar components \\object{HD\\,41004\\,A} and B. }  If confirmed, the present discovery represents the first detection of a short-period brown-dwarf companion around a M2 dwarf. { In particular, its estimated upper limit mass ($\\sim$25\\,M$_{\\mathrm{Jup}}$) puts it in the middle of the so-called brown-dwarf desert, a mass region (between $\\sim$20 and 40\\,M$_{\\mathrm{Jup}}$) for which (almost) no short period companions to solar-type F, G, K \\citep[][]{Hal00,Udr00b,Jor01} and M \\citep[][]{Mar89} dwarfs were found.} The fact that \\object{HD\\,41004\\,B} is a M dwarf makes us speculate that the formation of such systems is more likely for lower mass primaries, i.e. systems having mass ratios closer to unity \\citep[][]{Duq91}.  \\citet[][]{Arm02} proposed a model to explain the existence of the brown-dwarf desert. One interesting feature of their model is that it predicts that no brown-dwarf desert should be observed for the companions around the lowest mass dwarfs; they set an upper limit of 0.1-0.2\\,M$_{\\sun}$. Besides the fact that \\object{HD\\,41004\\,B} is slightly more massive than this limit, the present case is interestingly similar to the predictions.  Recently, \\citet[][]{Zuc02} discussed an interesting correlation between the planetary mass and the orbital period. Their analysis strongly suggests that for single stellar systems, there is a lack of ``high'' mass planetary companions in short period orbits. On the other hand, Zucker \\& Mazeh have found that for stars in multiple systems, this correlation is no longer present. In particular, they have pointed out that in these latter cases, there seems to be a negative correlation between ``planetary'' mass and orbital period (see their Fig.\\,3). It is very interesting to see that the companion to \\object{HD\\,41004\\,B} perfectly fits this trend.  A few works have tried to study the formation of planets (or low mass objects) around stars in multiple systems. \\citet[][]{Nel00} showed that the formation of a planet in a disk is unlikely for equal mass binary systems with a separation lower than $\\sim$50\\,AU. On the other hand, \\citet[][]{Bos98} suggested that the influence of a companion might trigger ``planetary'' formation by disk instability. Although no strong conclusions are possible at this moment, the fact that \\object{HD\\,41004\\,B} is in a double system with a separation that can be as low as $\\sim$20\\,AU is very interesting from the point of view of the formation of its companion.  Given the uncertainties in the shape of the observed bisector and in the models, we do not pretend to have a precise mass determination for the companion to \\object{HD\\,41004\\,B}. Although we believe we have obtained a good estimate, one of the main goals of this paper was to illustrate the importance of combining the radial-velocity data with the bisector analysis. In other words, one important lesson to be taken from the presented results is that the bisector analysis was crucial to correctly interpret the observations.  It is important to caution that this kind of situation can also happen for long period systems. For those, the signal is very unlikely to have an intrinsic stellar activity origin, but we cannot exclude that it might originate from the presence of a ``wobbling'' stellar companion. It is very easy to find a situation where an undetected companion, a few tens of an arcsec distant, can be ``contaminating'' our analysis. Together with a good knowledge of the target star environment, in such a situation the bisector analysis seems to provide an unique tool to point out the exact origin of the radial-velocity variations.  In this sense we have analyzed (or re-analyzed) the bisectors for all the stars with planets discovered in the context of the CORALIE planet search programme\\footnote{See obswww.unige.ch/$\\sim$udry/planet/planet.html}. We did not find any significant correlation between $V_r$ and BIS. We can thus remain confident that the presence of a planetary companion is the best way of explaining the radial-velocity variations in these systems.  \\begin{figure}[t] \\psfig{width=\\hsize,file=Fig13.eps} \\caption{Phase folded radial-velocity measurements of \\object{HD\\,41004\\,AB} obtained using two different CCF masks. As expected, the amplitudes $K$ obtained vary from mask to mask (see text). The vertical scales of the two plots is the same to facilitate a comparison.} \\label{fig_masks} \\end{figure}  Besides the bisector analysis, one other way of testing a radial-velocity variation for such cases may involve the measurement of the radial-velocity using different sets of lines, or different spectral regions. If the radial-velocity variation is due to the presence of a planet, we can expect that every spectral region/line will give us about the same velocity amplitude (but not necessarily the same $\\gamma$-velocity). The fact that for \\object{HD\\,41004\\,AB} the amplitude in radial velocity was different using the cross-correlation mask constructed specially for the bisector analysis \\citep[][]{Que01b} from the one obtained using the ``classical'' mask is very telling (see Sect.\\,\\ref{sec:bisector}). In Fig.\\,\\ref{fig_masks} we can further see two phase folded diagrams for the radial-velocities of \\object{HD\\,41004\\,AB} obtained using two different CCF masks. As we can see, if we use a mask specially constructed for the radial-velocity determination of M4 dwarfs \\citep[][]{Del98b}, i.e. a spectral type close to the one of \\object{HD\\,41004\\,B}, we obtain a much higher amplitude in radial-velocity than for the case of using a mask constructed for K0 dwarfs. This difference is expected since in such a case, the difference between the CCF's of \\object{HD\\,41004\\,A} and \\object{HD\\,41004\\,B} is much smaller (the M dwarf spectrum is enhanced relatively to the K dwarf), being thus the influence of the ``small'' CCF stronger. A similar situation (although not as strong) is seen when comparing the amplitudes obtained using the K0 and the F0 masks.  This kind of analysis can, in principle, serve as a test, if the study of the bisector is not possible. But we note that for radial-velocity variations induced by the presence of dark spots in the stellar photosphere, i.e. where the same phenomenon is affecting all spectral lines in about the same way, we do not expect a strong difference between the radial-velocity amplitudes obtained using different spectral lines.  Finally, it is interesting to say a few words about other possible ways of confirming the current detection. Given the short period of the brown dwarf around \\object{HD\\,41004\\,B}, the probability that we are able to observe a transit is quite high. In a first approximation, the magnitude variation expected in such a transit is around 10$^{-3}$ (for the A+B system), a value that does not seem too low. But a simpler way of confirming this case would pass by doing high-resolution spectroscopy (and velocity measurements) of \\object{HD\\,41004\\,B}. However, this is not an easy task, since the two components of \\object{HD\\,41004\\,AB} are separated by only 0.5\\arcsec. Unfortunately, there are no available high-resolution spectrographs attached to an Adaptative Optics system in the southern hemisphere capable of accomplishing this task. The use of the Hubble Space Telescope (HST) might represent a solution. Else, the solution may pass by using high-resolution near-IR spectroscopy, since at those wavelengths the flux of \\object{HD\\,41004\\,B} is much closer to the one from \\object{HD\\,41004\\,A}."
    },
    "0206/astro-ph0206160.txt": {
        "abstract": "A {\\sl Chandra X-Ray Observatory} ACIS-S imaging observation is used to study the population of X-ray sources in the nearby Sab galaxy M81 (NGC 3031). % A total of 177 sources are detected with 124 located within the $D_{25}$ isophote to a limiting X-ray luminosity of $\\sim$$3\\times 10^{36}$ \\ergl. % Source positions, count rates, luminosities in the 0.3~--~8.0~keV band, limiting optical magnitudes, and potential counterpart identifications are tabulated. Spectral and timing analysis of the 36 brightest sources are reported including the low-luminosity active galactic nucleus, SN~1993J, and the \\ein-discovered ultra-luminous X-ray source X6. The nucleus accounts for $\\sim$86\\%, or $5\\times$$10^{40}$~\\ergl, of the total X-ray emission from M81. Its spectrum is well-fit by an absorbed power law with photon index 1.98$\\pm$0.08 consistent with previous observations (average index 1.9). SN~1993J has softened and faded since its discovery. At an age of 2594 days, SN~1993J displayed a complex thermal spectrum from a reverse shock rich in Fe~L and highly-ionized Mg, Si, and S but lacking O. A hard X-ray component, emitted by a forward shock, is also present. X6 is spatially-coincident with a stellar object with optical brightness and colors consistent with an O9~--~B1 main sequence star. It is also coincident with a weak radio source with a flux density of $\\sim$95~$\\mu$Jy at $\\lambda=3.6$~cm. The continuum-dominated X-ray spectrum of X6 is most closely reproduced by a blackbody disk model suggesting the X-ray source is an $\\sim$18~\\msun\\ object accreting at nearly its Eddington limit.  The non-nuclear point source population of M81 accounts for 88\\% of the non-nuclear X-ray luminosity of $8.1\\times10^{39}$~\\ergl. The remaining (unresolved) X-ray emission is confined within $\\sim$2~kpc of the galactic center. The spatial distribution of this emission and of the resolved X-ray bulge sources closely follows that of the bulge optical light. In particular, there is no evidence for an X-ray signature accompanying the filamentary H$\\alpha$ or excess UV emission seen in the central \\LA1.0~kpc of the galaxy. The shape of the luminosity function of the bulge sources is a power law with a break at $\\sim4\\times10^{37}$~\\ergl; suggesting the presence of an aging ($\\sim$400~Myr) population of low-mass X-ray binaries. Extrapolating this luminosity function to lower luminosities  accounts for only $\\sim$10\\% of the unresolved X-ray emission. Spectroscopically, the unresolved emission can be represented as a combination of soft, $kT$$\\sim$0.3~keV, optically-thin plasma emission and of a $\\Gamma=1.6$ power law. The unresolved bulge X-ray emission is therefore most likely a combination of hot gas and of one or more large and distinct populations of low-luminosity X-ray sources confined in the gravitational potential and tracing the old population of bulge stars. % The distribution of disk sources shows a remarkably strong correlation with spiral arms with the brightest disk sources located closest to spiral arms. The luminosity function of sources near the spiral arms is a pure power law (slope $-0.48\\pm0.03$) while that of sources further away exhibits a break or cut-off in the power law distribution with no high-luminosity members. This is interpreted as a natural consequence of the passage of spiral density waves that leave the brightest (when averaged over their lifetimes) and shortest-lived X-ray sources immediately downstream of the spiral arms. Consistent with model predictions, we conclude that the shapes of the X-ray luminosity functions of the different galactic components of M81 are most likely governed by the birth rates and lifespans of their constituent X-ray source populations and that the luminosity functions can be used as a measure of the star formation histories of their environments. ",
        "introduction": "Systematic investigations of the X-ray properties of normal galaxies began in earnest with the \\ein\\ observatory over two decades ago. The picture that emerged for spiral galaxies (see the early reviews by Long \\& van~Speybroeck 1983; Helfand 1984; and Fabbiano 1989) is that the bulk of the X-ray emission takes place in two distinct physical environments: the star-forming disks of late-type spiral and irregular galaxies and among the old stellar population in dense globular clusters and compact bulges at the centers of early-type spiral galaxies. % In addition to these trends along the Hubble sequence, variations were sometimes found among spiral galaxies of similar morphological type suggesting a dependence on star formation histories. In particular, the brightest X-ray emissions are associated with starbursts in merging and interacting galaxies (David, Jones, \\& Forman 1992). Thus, by tracing the endpoints of stellar evolution, the X-ray source populations of external galaxies provide important clues to the physical nature and evolutionary history of their hosts.  The contemporary view for spiral galaxies is rapidly being refined following the launch of the \\cha\\ and \\xmm\\ X-ray Observatories. % Moderately-deep images reveal point sources to limiting X-ray luminosities of order $10^{37}$~\\ergl\\ in galaxies out to Virgo cluster distances. While this samples only the high luminosity end of the distribution of X-ray sources, of order 100 sources are routinely detected in normal galaxies similar to our own. Reliable spectral analysis is usually limited to an even smaller subset of the brightest individual sources. Nevertheless, using probabilistic methods, the observed sample of X-ray sources can help us understand current-epoch galaxy evolution in its broader context.   A formal expression of the relationship between the star formation history of local galaxies and their observed X-ray source populations has recently been put forth by Wu (2001; see also Wu \\etal\\ 2002a,b; Kilgard \\etal\\ 2002; Dalton \\& Sarazin 1995). There it was shown that the basic shape of the observed X-ray luminosity function is governed simply by the birth and death rates of the source population under the assumption that the more luminous X-ray sources are shorter-lived. Thus, in the absence of ongoing star formation, the luminosity function will develop a cutoff at high luminosity that evolves toward lower luminosity. Conversely, if the population of X-ray sources is replenished through star formation processes, such as is found in spiral arms, then a power law shaped luminosity function can be sustained.  Certain complications arise when applying this basic interpretation to X-ray populations in individual galaxies (Wu \\etal\\ 2002a,b). Among these are the presence of different classes of X-ray sources, such as supernova remnants and accreting compact objects, which evolve on differing timescales; alternative source-formation mechanisms uncorrelated with stellar evolution such as binary captures in globular clusters; and non-steady or luminosity-limited emission characteristics such as those associated with X-ray transients and novae and in Eddington-limited neutron star binaries, respectively. % Therefore, only when specific counterparts to individual X-ray sources are identified and their multiwavelength properties assessed can the full power of the hypothesis of Wu \\etal\\ (2002a,b) be applied to address the nature and evolution of X-ray sources in different environments.  The nearby Sab galaxy M81 (NGC 3031) is ideal for such a study in that it contains both a strong two-arm grand-design spiral pattern and a well-defined circumnuclear bulge. The distance to M81, 3.6~Mpc, has been well-established from Cepheid measurements (Freeman \\etal\\ 1994) which are in good agreement with other distance estimates (Ferrarese \\etal\\ 2000). Populations of several classes of objects in M81 have been investigated and catalogued including globular clusters (Perelmuter \\& Racine 1995; Chandar, Ford, \\& Tsvetanov 2001), supernova remnants (Matonick \\& Fesen 1997), H~II regions (Hodge \\& Kennicutt 1983; Petit, Sivan \\& Karachentsev 1988), and stars and star clusters (Zickgraf \\& Humphreys 1991; Ivanova 1992; Sholukhova \\etal\\ 1998). In addition, the plane of the galaxy is oriented 32$^{\\circ}$ from face-on allowing detailed mapping of the velocity field (Goad 1976; Rots \\& Shane 1975; Adler \\& Westphal 1996) for dynamical studies and testing spiral density wave models (Visser 1980; Roberts \\& Hausman 1984).  The center of M81 contains a compact radio core (Bartel \\etal\\ 1982) surrounded by a region of enhanced far-infrared (Rice 1993; Davidge \\& Courteau 1999), H$\\alpha$ (Devereux, Jacoby \\& Ciardullo 1995), and ultraviolet (Hill \\etal\\ 1992; Reichen \\etal\\ 1994) emission extending to $\\sim50\\arcsec$  ($\\sim$900 parsecs). This emission probably comes from an old population of hot, low-mass stars rather than from young massive stars (O'Connell \\etal\\ 1992; Devereux, Ford \\& Jacoby 1997).  In contrast to the bulge, HI velocity contours show a sharp discontinuity beyond the bulge identified as a spiral velocity shock (Visser 1980). Downstream of this shock are regions of star formation in the spiral arms. The distributions of these components are consistent (Kaufman \\etal\\ 1989) with density wave models predicting a broad spiral density enhancement (e.g., Roberts \\& Hausman 1984).  Beyond the visible disk of M81 is an envelope of neutral hydrogen (Roberts 1972) enclosing M81 and nearby group members. A bridge of gas, a relic of tidal interaction (Cottrell 1977), connects M81 and the starburst galaxy M82.   The hypothesis of Wu (2001) and Wu \\etal\\ (2002a,b) was motivated in large part by the initial results from our \\cha\\ observation of M81 presented in Tennant \\etal\\ (2001). There it was shown that the X-ray luminosity function of the population of bulge sources displays a break at $\\sim$4$\\times$$10^{37}$ \\ergl\\ similar to that observed in M31 (e.g., Shirey \\etal\\ 2001). The X-ray luminosity function of the disk sources, on the other hand, follows a single power law slope over three decades in flux. This is what is expected if an impulsive episode of star formation occurred in the bulge in the past, ostensibly during an encounter between M81 and one of its companion galaxies, while continuous star formation in the disk is being driven by the passage of spiral density waves.  Here we build upon the earlier work of Tennant \\etal\\ (2001). After presenting detailed information on the individual X-ray sources in \\S\\ref{s:discrete_src} and in-depth analysis of the brightest objects in \\S\\ref{s:bright_src}, the properties of the bulge (\\S\\ref{s:bulge}) and disk (\\S\\ref{s:disk}) regions are addressed separately then discussed (\\S\\ref{s:discussion}) within the common framework of galaxy evolution.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{s:discussion} \\subsection{The Discrete X-ray Source Population}  The X-ray spectra of the brightest sources in the M81 field are predominantly moderately-absorbed power laws with photon index $\\Gamma\\sim1.5$. A similar spectral shape reproduces the combined bulge resolved-source spectrum while the combined disk source spectrum (for sources on S3) requires an additional weak thermal component. A power law is indicative of accreting X-ray binaries (XRBs) and a large population of bright power law sources is consistent with surveys of our Galaxy and the Local Group where the bright X-ray source population is dominated by low-mass XRBs ({\\sl e.g.}, Grimm \\etal\\ 2001).  The scarcity of X-ray sources detected in the radio, or correlated with optically-selected SNRs, or exhibiting a strong thermal X-ray spectrum implies an insignificant number of the bright X-ray sources in M81 are SNRs. Supernova remnants are common in the Magellanic Clouds and in the solar neighborhood but are relatively less luminous and short-lived compared to typical XRBs. Perhaps some of the weaker disk sources are SNRs and account for the thermal emission present in the combined disk source spectrum. Multi-wavelength observations of optically-identified extra-galactic SNRs have revealed that these sources are typically very weak X-ray and radio emitters (e.g. Pannuti \\etal\\ 2000, Lacey \\& Duric 2001, Pannuti \\etal\\ 2002).  Four X-ray sources are coincident with known M81 globular clusters 3 other X-ray sources are coincident with optically-bright objects with colors consistent with globular clusters and two other X-ray sources have candidate optical counterparts that appear extended in \\hst\\ images. Approximately 10\\% of Galactic globular clusters contain X-ray sources. In the case of M81, two surveys of the globular cluster population report a total of 139 globular clusters including 98 within the \\cxo\\ field of view. Thus, 4\\% to 7\\% are coincident with X-ray sources. DiStefano \\etal\\ (2002) find 25\\% of confirmed clusters in their M31 field of view contain X-ray sources and 10\\% of all globular cluster candidates have X-ray sources. They also report that most of the luminous M31 X-ray sources are in globular clusters. In contrast, only one of the 11 brightest M81 X-ray sources, \\ein\\ source X7, is coincident with a globular cluster.  In general, there are remarkably few counterparts to the resolved X-ray sources identified in our assessment of the extensive literature and available archival images of M81. If the majority of the resolved sources are XRBs, then they have companion stars and accretion disks that may be detectable in optical light. The \\hst\\ images of M81 approach a limiting magnitude of $V\\sim27$~mag or $M_V\\sim-0.5$~mag. Thus only O and B main sequence or giant companions or highly-luminous accretion disks from XRBs in active states will be detectable by \\hst. Later-type companion stars will not appear in the optical data.  Low-mass XRBs are common in galaxies because they are long-lived and slowly-evolving. The last encounter of M81 and its companion galaxy M82 occurred some 500 Myr ago based on the study of the ages of young star clusters in M82 (de~Grijs, O'Connell, \\& Gallagher 2001). If the onset of the last major star formation episode in the bulge of M81 was triggered by this encounter, then the most massive members of the current population of main-sequence stars have masses \\LA2.5~\\msun\\ (Maeder \\& Meynet 1988). If these constitute the population of companion stars in the currently-active XRBs, then the resolved bulge sources are mostly low-mass XRBs. They will not have detectable optical counterparts because they do not have O, B, or giant star companions (and because of the bright, amorphous optical background of the bulge). In contrast, on-going star formation along the spiral arms should produce a population of high-mass XRBs with massive O and B star companions. This environment is, however, also the location of obscuring atomic and molecular gas. The few correlations with HII regions suggests some of the X-ray sources are located in star forming regions that may be populated by massive stars. The high percentage of \\hst\\ potential counterparts identified in the disk relative to the bulge also suggests an abundance of early-type stars in the vicinity of the disk sources and, potentially, a preference for high-mass XRB systems in this environment. This is consistent with the distribution of XRBs in our Galaxy where high-mass XRBs are concentrated towards the Galactic plane and along spiral arms while low-mass XRBs show a concentration towards the Galactic center (Grimm \\etal\\ 2001).  \\subsection{The Brightest M81 Sources}  In-depth analysis of three of the 4 brightest sources in the M81 field was presented in \\S\\ref{s:bright_src} and of the third-brightest source in Swartz \\etal\\ (2002). Interestingly, all three of the brightest non-nuclear source are far from typical XRBs as seen in our Galaxy. SN~1993J is a supernova, \\ein\\ source X6 is a rare ultra-luminous X-ray source with possible optical and radio counterparts, and source number 132 is an exceptionally-bright and hot supersoft source candidate (Swartz \\etal\\ 2002).  \\subsubsection{SN 1993J}  SN~1993J appears to be evolving as expected based on the standard CSM interaction model of Chevalier (1982; Fransson \\etal\\ 1996) though a complete picture incorporating models of the exploding star and its pre-supernova environment awaits detailed numerical calculation. The X-ray properties of SN~1993J reported here provide an important constraint on any future theoretical investigations because the X-ray light curve is declining steadily, even at $\\sim$7~yr, whereas the most-detailed numerical simulations to date (Suzuki \\& Nomoto 1995) predicted the light curve would drop precipitously long ago unless the CSM were clumpy. A clumpy CSM would produce a varying light curve with episodes of high X-ray flux occuring whenever clumps are overtaken by the outgoing shock wave (Chugai 1993). In this scenario, the CSM consists of a rarified wind embedded with relatively dense clouds and the X-rays emanate from the shocked gas of the clouds with little or no reverse shock emission. This is not what is observed spectroscopically. The spectrum of SN~1993J is best modeled with a combination of thermal emission from a reverse shock and a hard component from a forward shock.  \\subsubsection{Einstein Source X6}  The multi-wavelength properties of \\ein\\ source X6 are intriguing. The X-ray spectrum of the source is best-fit with a disk blackbody model. In this model, the X-rays come from the inner portions of an accretion disk surrounding a compact object. The inferred mass of the central object is $\\sim$18~\\msun\\ assuming the innermost disk radius derived from the model corresponds to the last stable Keplerian orbit of a non-spinning black hole. The X-ray-model-derived bolometric luminosity is near the Eddington limit for an object of this mass. The X-ray flux from X6 has been persistent throughout the $>$20~yr of observation.  X6 is located within a 5$\\arcsec$-diameter ($\\sim$90~pc) SNR candidate according to Matonick \\& Fesen (1997) based on a high [SII]/H$\\alpha$ ratio indicative of collisional excitation in the cooling region behind a SNR shock. There are no emission lines present in the X-ray spectrum of X6 and an optically-thin thermal plasma model is a notably poorer fit to the X-ray data. The X-ray morphology of X6 is that of a point source with no evidence for extension. Thus, no X-ray evidence, besides a steady flux, supports the conjecture that X6 is a SNR.  A weak radio source is  present at the location of X6. Synchrotron emission is observed at radio wavelengths from relativistic electrons accelerated in SNR shocks and in jets emanating from some (Galactic) XRBs. A radio (or optical) light curve of the source at the location of X6 has not yet been constructed. The radio source was present at 3.6~cm in late 1994 but not seen in a 6~cm image taken in late 1999. Analysis of other radio images is in progress. If the radio source proves to be variable, then it is not from a SNR. The observed radio flux density of $\\sim$95~$\\mu$Jy is typical of, for example, Magellanic Cloud SNRs (Filipovic \\etal\\ 1998) after accounting for the disparate distances. In comparison, radio jets associated with Galactic XRBs are weaker except during extreme outbursts.  There is also an optical point source coincident with X6. If associated with X6, the optical emission may either be from a moderately-massive, O9~--~B1, companion, which may be a Be star, or from the accretion disk itself but does not come from an extended source at \\hst\\ resolution.  X6 can be compared to well-studied nearby XRBs. An example of a high-mass system with a massive compact accretor is Cyg~X-1 (e.g., van~Paradijs 1995). The optical counterpart to Cyg~X-1 is a O9.7 supergiant with colors similar to those of the X6 counterpart but with the higher optical luminosity of a supergiant compared to a main-sequence star. The putative black hole in Cyg~X-1 exceeds $7$~\\msun\\ and is most probably $\\sim$16~\\msun. Cyg~X-1 is a persistent X-ray source as is X6. It displays the charactersitic high-soft and low-hard states typical of black hole XRBs (Tanaka \\& Lewin 1995) and does not exceed an X-ray luminosity of $\\sim2\\times10^{38}$~\\ergl. Cyg~X-1 is radio-bright during its low-hard state with a flux of $\\sim$15~mJy or $0.007$~$\\mu$Jy if it were placed at the distance of M81. Scaling this value upward by the ratio of the X-ray luminosities of X6 to Cyg~X-1 in its low-hard state ($\\sim$700) results in a radio flux density of only $6$~$\\mu$Jy which would not be detectable.  An example of a low-mass system with a massive compact object and strong radio emission is the microquasar GRS~1915+105 (e.g., Mirabel \\& Rodr\\`{i}guez 1999). This system is a rapidly variable X-ray and radio transient reaching a peak X-ray luminosity of $\\sim1.5\\times10^{39}$~\\ergl\\ in its high-soft state, comparable to X6, and an average luminosity of $\\sim3.7\\times10^{38}$~\\ergl. High extinction along the line of sight to GRS~1915+105 obscures the optical counterpart and accretion disk. Near-infrared spectroscopy (Greiner \\etal\\ 2001), however, shows the companion to be a K~--~M main-sequence star and, along with the orbital period, constrained the compact object mass to be $14\\pm4$~\\msun. GRS~1915+105 is a strong radio emitter. Scaling to the distance of M81 and to the X-ray luminosity of X6 (a factor of $\\sim$13 when GRS~1915+105 is in its hard state) results in a radio flux density of about one-half the X6 value.  Thus, while monitoring at many wavelengths is required before any definitive statment can be made, it is intriguing to consider X6 may be an X-ray- and radio-bright member of the class of microquasars (see Mirabel \\& Rodr\\`{i}guez 1999 for a review) consisting of an accreting black hole with a radio-bright jet but with unusually-steady X-ray flux.  \\subsubsection{The M81 Nucleus}  The X-ray properties of the nucleus of M81 are difficult to deduce from the present dataset because of severe pileup. A relatively weak spectrum extracted from the readout trail was analyzed and found to be a power law of photon index $\\Gamma = 1.98\\pm0.08$, consistent with numerous previous X-ray studies. The presence of Fe~K$\\alpha$ emission could not be confirmed because of the lack of counts above $\\sim$5~keV. Variability of the source also could not be assessed. However, it was shown, with the aid of the high angular resolution of the \\cxo\\ image, that the contribution to the nuclear spectrum from thermal emission is small or non-existent. Any thermal X-ray component present in the region is consistent with an extrapolation of the unresolved bulge emission observed surrounding the nucleus and extending over an $\\sim$4~kpc diameter region.  \\subsection{The M81 Bulge}  In addition to 53 X-ray sources resolved in the \\cxo\\ image, the bulge of M81 emits $\\sim10^{39}$~\\ergl\\ in unresolved emission. This is $\\sim$12\\% of the total non-nuclear emission from the entire galaxy and is distributed over an $\\sim$2$\\arcmin$-radius region centered on the nucleus. Both the resolved sources and unresolved emission trace the optical light from the old population of bulge stars.  If the unresolved emission is also produced by stellar systems, then they are systems distinct from the resolved sources because extrapolation of the luminosity function of the resolved sources contributes $<$10\\% of the unresolved emission. The possible X-ray-luminous stellar systems below the detection limit are massive OB stars, Be XRBs, CVs, RS CVn stars, and, at a lower luminosity, late-type stars. However, individual late-type stars have X-ray luminosities only of order a few $10^{27}$ to a few $10^{28}$~\\ergl, requiring some $10^{12}$ stars to produce the unresolved emission.  Massive OB stars with colliding winds can be strong X-ray emitters but are rare. None are found in the bulge of M81 (Devereux, Ford, \\& Jacoby 1997). Be XRBs are young high-mass systems and the Be companion star is optically bright. They are therefore also  unlikely to be abundant in the galactic bulge.  CVs are short-period (typically $<$1 day) binaries consisting of a white dwarf and a late-type low-mass companion (Warner 1995). They are numerous and are long-lived. The magnetic CVs, with a magnetic white dwarf, are known to have X-ray luminosities as high as $\\sim10^{32}$~\\ergl. The space density of magnetic CVs in the solar neighborhood is $\\sim 10^{-6}$~pc$^{-3}$ (Warner 1995) while the stellar density is about 0.7~\\msun~pc$^{-3}$ (Allen 1973). This implies the density of magnetic CVs is $\\sim 10^{-5}$~\\msun$^{-1}$. If M81 has a similar space density of magnetic CVs, then there will be roughly $10^5$ CV systems in the bulge of M81. If about 10\\% are active (a rough estimate based on the properties of the local systems), then only $\\sim10^{36}$~\\ergl\\ of the unresolved bulge X-ray emission can come from CVs.  RS~CVn systems, composed of chromospherically active G or K stars with late-type main sequence or subgiant companions, also have high X-ray luminosities. Typical X-ray luminosities of RS~CVn systems range from $\\sim10^{29}$~\\ergl\\ to $\\sim3\\times10^{31}$~\\ergl\\ (Rosner, Golub, \\& Vaiana 1985). Thus some $10^7$ to $10^{10}$ RS~CVn systems are required to produce the unresolved bulge emission. If all stars in the bulge are $\\sim$1~\\msun\\ and half are in binary systems, then there are $\\sim$$10^{10}$ binary systems in the bulge. Assuming about 20\\% of these systems become RS~CVns and that G-K stars spend only a few percent of their lifetimes in their giant stage, an uncomfortably large fraction must currently be in an RS~CVn phase.  Individually, therefore, none of the these stellar systems can readily account for the observed unresolved X-ray emission from the M81 bulge. % If, instead, some portion of the unresolved bulge X-ray emission is from hot diffuse gas, as suggested by its spectral distinction from the simple power law shape of the resolved sources, then only a small fraction ($\\sim$0.02\\%) of the total bulge mass is needed to account for the observed emission. The X-ray emission, however, does not appear filamentary like the H$\\alpha$ emission does (Devereux \\etal\\ 1995). A filamentary morphology would be expected if the emission is from ionization by shocks. A source for producing shocks is also not obvious. Devereux \\etal\\ (1997) suggest shocks originating from nuclear activity can account for the wispy ``nuclear spiral'' of H$\\alpha$ emission confined to the central $\\sim$1$\\arcmin$  but the unresolved X-ray emission is rather smoothly extended over a region of 2$\\arcmin$ radius. While ionizing radiation from hot evolved post-AGB stars may produce the observed UV excess in the core of M81 (O'Connell \\etal\\ 1992, Devereux \\etal\\ 1995), these stars do not produce adequate ionizing radition in the \\cxo\\ energy band to account for the X-ray emission (Binette \\etal\\ 1994).  \\subsection{The M81 Disk and Spiral Arms}  One of the most spectacular features of M81 is its grand design spiral arm structure. The spiral arms trace the location of recent star-forming activity induced by the passage of spiral density waves. Applications of classical density wave models to M81 (e.g., Visser 1980) predict that material travels faster than the spiral pattern, entering an arm on the inside ``upstream'' edge. Stars forming at the spiral shock front travel at the local circular velocity of galactic rotation so that the youngest stars would be immediately downstream of the shock or toward the outside edge of the arm. The most massive stars are the quickest to evolve. They end their lives in core-collapse SN explosions leaving behind a neutron star or, perhaps, a black hole remant. Supernova explosions produce X-ray emitting SNRs and compact stars in binaries may become XRBs. Thus, the spiral arms are not only the site of star formation but also a stellar graveyard and the birthplace of X-ray sources.   The brightest X-ray sources in the disk of M81 correlate spatially with the spiral arms. Accepting that the majority of the resolved sources are XRBs and that the X-ray flux is generally proportional to the mass accretion rate, then the brightest XRBs are young high-mass XRBs with high mass-transfer rates. The onset of mass-transfer in these systems, and hence of the X-ray-bright phase, can begin immediately following the formation of the compact object because of the strong stellar wind from the massive companion. This is in contrast to the low-mass XRBs in which mass transfer begins only after the companion star evolves to a (sub)giant stage or when the binary orbit has decayed sufficiently so that Roche lobe overflow can begin. Thus, the young high-mass systems become X-ray emitters while still within the spiral arm region of their origin. For this reason the luminosity distribution of the young XRBs in the spiral arms are expected to differ from the distribution of the older XRBs elsewhere in the galaxy. In particular, it will not show the characteristic luminosity break induced by aging of the XRB population as predicted by Wu (2001) and Wu \\etal\\ (2002a,b).  Core-collapse supernova only come from stars more massive than $\\sim$8--10~\\msun\\ and are X-ray bright SNRs only for a relatively short time. They, too, should be found preferentially near their place of origin, the spiral arms.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   The observed M81 X-ray luminosity functions reported by Tennant \\etal\\ (2001) prompted Wu (2001) and Wu \\etal\\ (2002a,b) to consider the physical underpinnings that give rise to a cutoff in the luminosity function for the bulge sources and to the absence of this feature in the disk population. Wu (2001) showed that the shape of the luminosity function is governed to first order simply by the birth rates and (X-ray active) lifespans of the XRBs that dominate the luminosity function and hence is a measure of the star formation history of the local environment and of galaxy evolution in the broader context. Further investigation (Wu \\etal\\ 2002a,b) reveals that several complicating issues must be considered before this hypothesis can be rigorously applied. Some of these issues have been addressed in the present work.  The first issue is the presence of a population of SNRs. The onset and duration of the X-ray active phase of XRBs depends mainly on the donor star mass and its consequent evolutionary path while the X-ray luminosity depends on the accretion rate. This is fundamentally different than the X-ray evolution of SNRs. Here we have shown, however, that SNRs are not important contributors to the total X-ray source population in M81 with the exception of SN~1993J, the fourth-brightest source in the M81 field at the time of observation.  Another issue is the occurrence of XRBs in globular clusters. If capture processes govern the formation of XRBs in globular clusters, as seems likely to account for the excess of XRBs in these environments, then the characteristic lifetime of the XRB is not just correlated with the nuclear or orbital evolution timescales of the system but is also a function of the encounter frequency. However, only a few percent of the X-ray sources in M81 appear to be in globular clusters. As with Galactic globular cluster XRBs (Verbunt \\& van~den~Heuvel 1995), the impact on the luminosity function is further minimized by the fact that the globular cluster XRBs in M81 are not among the brightest X-ray sources. Again, the exception is \\ein\\ source X7, the fifth-brightest source in the field.  A third factor with potential impact to the basic hypothesis of Wu (2001) is the presence of XRBs with a nuclear-burning white dwarf accretor, i.e., members of the class of supersoft sources (e.g., Kahabka \\& van~den~Heuvel 1997). While the lifespans of these objects depend on the companion mass and mass-transfer rates as with other XRBs, only a narrow range of mass-transfer rates result in {\\sl steady} nuclear burning. Wu \\etal\\ (2002a,b) argue, therefore, a narrow X-ray luminosity range for this source population and a sharp decline in the number of sources with luminosities above the Eddington limit for a Chandrasekar-mass accretor. Swartz \\etal\\ (2002) found 9 supersoft source candidates in the M81 field. Six of these are relatively weak sources with luminosities in a narrow range around $\\sim10^{37}$~\\ergl. The three brightest candidates, however, radiate at or above the Eddington limit depending on the adopted spectral model with one source (the third-brightest in the entire field) approaching $\\sim10^{39}$~\\ergl. Thus the most important contribution to the luminosity function from supersoft sources is at the high luminosity end and is dominated by one bright source. The effect of the weaker supersoft sources is obscured by the large number of other sources contributing at low luminosities.  \\begin{center} \\includegraphics[angle=270,width=\\columnwidth]{f17.eps} \\vspace{10pt} \\figcaption{Observed X-ray luminosity function including all 123 non-nuclear X-ray sources detected within the $D_{25}$ isophote of M81 ({\\em heavy solid line}). Also shown are the luminosity functions of the SNRs ({\\em dotted}), of X-rays sources spatially-coincident with globular clusters {\\em dashed}), and of the supersoft sources ({\\em dot-dashed}). The luminosity function with these three populations and the Ultra-Luminous X-ray source X6 omitted is shown as a thin solid line. The symbol ($\\star$) marks the division between bright sources for which spectral analysis has been reported in this work and of weak sources for which no spectral fits were made. \\label{f:f17}} \\end{center}  The luminosity function for all non-nuclear sources detected within the $D_{25}$ isophote of M81 is shown in Figure~\\ref{f:f17}. Also shown are the luminosity functions of the supersoft sources, of X-ray sources spatially-coincident with globular clusters, and with SNRs. As shown in Figure~\\ref{f:f17}, these three populations all have relatively flat power-law luminosity functions and they affect only the bright end of the overall luminosity distribution. XRBs, the dominant population of X-ray sources in M81, however, have a steep luminosity function and hence determine the overall shape of the luminosity functions, especially at the faint ends. The break at the luminosity of $\\sim 4 \\times 10^{37}$~erg~s$^{-1}$ that we have found (see also Tennant \\etal\\ 2001) is therefore a characteristic imprint of the XRBs. We have argued that the formation of such a break is due to the age of a population of XRBs which were born at a star-burst episode in the recent past (Wu 2001, Wu \\etal\\ 2002a,b). This break is distinguishable from another possible break, expected to occur at $\\sim 2 \\times 10^{38}$~erg~s${-1}$, the Eddington luminosity of a 1.5~\\msun\\ accreting object. The latter is attributed (Sarazin, Irwin, \\& Bregman 2001) to the presence of a population of neutron stars which accrete at rates close to the Eddington limit. Whether or not this break is visible in a given population depends on the relative proportion of neutron-star XRBs and black-hole XRBs. It also requires that the host galaxies (e.g. giant ellipticals) have a sufficiently large X-ray source population that the break becomes statistically significant. Nevertheless, we do see hints of this break in the luminosity function of the X-ray sources in M81, when we remove the SNRs, globular clusters XRBs and the supersoft sources.  In summary, most of the factors complicating the simple birth-death model are unimportant for M81. Nevertheless, careful examination of the brightest sources is warranted because they have the largest influence on the luminosity function and yet are certainly not typical of the dominant class of X-ray sources, the XRBs.  \\vspace{0.125in} % acknowledgements % We thank L. Townsley for applying her CTI-corrector algorithm to the \\cha\\ data and for providing response matrices. We thank J. Davis for discussions and his independent analysis of the spectrum of the nucleus. We are grateful to N. Bartel and M. Beitenholtz for sharing their radio data. T.~G.~P. acknowledges a travel grant from the NRAO to reduce the radio data and he is very grateful to M. Rupen for guidance during the reduction process. K.~W. thanks M. Weisskopf for  providing support for visits to MSFC. Support for this research was provided in part by NASA/\\cha\\ grants GO0-1058X and AR2-3008X to D.~A.~S."
    },
    "0206/astro-ph0206025_arXiv.txt": {
        "abstract": "Models of thermal emission of neutron stars, presumably formed in their atmospheres, are needed to infer the surface temperatures, magnetic fields, chemical composition, and neutron star masses and radii from the observational data. This information, supplemented with  model equations of state and neutron star cooling models, is expected to move us further in understanding the fundamental properties of the superdense matter in the neutron star interiors. The neutron star atmospheres are very different from those of usual stars due to the immense gravity and huge magnetic fields. In this presentation we review the current status of the neutron star atmosphere modeling and present most important results. ",
        "introduction": "A systematic study of X-ray emission from isolated neutron stars (NSs), including radio pulsars, has started after the launch of the {\\sl Einstein} and {\\sl EXOSAT} space observatories. These studies have shown that, generally, there are two different components of the NS X-ray emission --- thermal and nonthermal. The nonthermal component with a power-law spectrum, observed from many radio pulsars, is believed to originate from the pulsar's magnetosphere, while the thermal component is emitted from the NS surface layers (atmospheres). The thermal radiation is particularly interesting because it can provide important information about the NS: its surface temperature, magnetic field, and chemical composition, as well as the NS radius and mass. Measuring these parameters for a sample of NSs is necessary for studying the thermal evolution of NSs and constraining the equation of state and composition of the superdense matter in the NS interiors (see the review by Yakovlev et al. in these Proceedings).  Thermal X-ray emission from the NS surface had been discussed by Chiu \\& Salpeter and Tsuruta in 1964, before NSs were discovered, and well before their thermal emission was detected with {\\sl Einstein} and {\\sl EXOSAT} (e.g., Cheng \\& Helfand 1983; Brinkmann \\& \\\"Ogelman 1987; C\\'ordova et al.~1989; Kellet at al.~1987). Many new observational results on the NS thermal radiation were obtained in 1990' with the {\\sl ROSAT}, {\\sl ASCA}, and {\\sl EUVE} satellites (see Becker \\& Pavlov 2002 for a review). Thermal radiation from a few NSs was also detected in the optical-UV energy range with the {\\sl Hubble Space Telescope} (e.~g., Pavlov, Stringfellow \\& C\\'ordova 1996; Walter \\& Matthews 1997). Currently operating X-ray observatories, {\\sl Chandra} and {\\sl XMM}-Newton, are providing new excellent data on X-rays from NSs (see the contributions by Weisskopf, Becker, and Pavlov, Zavlin \\& Sanwal in these Proceedings).  To interpret these observations, one needs reliable models for the NS thermal radiation. The fact that the spectrum of radiation emergent from a NS atmosphere can be very different from a blackbody spectrum, and the angular distribution can be very far from isotropic, particular in a strong magnetic field, has been recognized long ago. For instance, Pavlov \\& Shibanov (1978) calculated spectra and angular distributions of radiation from a strongly magnetized NS atmosphere assuming the source function grows inward linearly in the emitting layers. First self-consistent models for NS atmospheres were developed by Romani (1987), for low magnetic fields, and Shibanov et al.\\ (1992), for strong magnetic fields. Since then, various aspects of the NS atmosphere modeling have been investigated in many papers.     Below we will overview the current status of this field and summarize some important results. ",
        "conclusions": "In the recent decade, substantial progress has been made in modeling atmospheres of isolated NSs. Best investigated cases are nonmagnetic atmospheres and fully-ionized light-element atmospheres with strong magnetic fields. The atmosphere models have been applied to the interpretation of thermal emission from NSs of different types. For instance, Pavlov \\& Zavlin (1997), Zavlin \\& Pavlov (1998), and Zavlin et al.\\ (2002) analyzed the X-ray emission from the millisecond pulsar J0437--4715 using the nonmagnetic light-element atmosphere models; Pavlov et al.\\ (2001) applied the magnetic hydrogen models to the analysis of radiation from the Vela pulsar; Zavlin, Pavlov \\& Tr\\\"umper (1998) and Zavlin, Tr\\\"umper \\& Pavlov (1999) used the magnetic light-element models to interpret the thermal emission from NSs in the supernova remnants PKS 1209--51/52 and Puppis~A. More examples are presented in the contribution by Pavlov et al.\\ in this volume.  However, a number of problems in the atmosphere modeling remains to be solved. First of all, investigations of the structure of various atoms, molecules, and molecular chains in strong magnetic fields, as well as radiative transitions in these species (Pavlov 1998), are necessary to construct magnetic atmosphere models of different chemical compositions. First efforts in this direction are being undertaken (Mori \\& Hailey 2002). Particularly interesting are the (virtually unknown) radiative properties of matter in superstrong magnetic fields, $B\\ga 10^{14}$ G, apparently found in anomalous X-ray pulsars and soft gamma-ray repeaters. Further work is needed on radiative properties of nonideal plasmas and condensed matter. These investigations will eventually result in more advanced models for thermal radiation of isolated NS, needed for the interpretation of high-quality observations of these objects with the {\\sl Chandra} and {\\sl XMM}-Newton X-ray observatories."
    },
    "0206/hep-th0206138_arXiv.txt": {
        "abstract": "The ekpyrotic and cyclic universe scenarios have revived the idea that the density perturbations apparent in today's universe could have been generated in a `pre-singularity' epoch before the big bang. These scenarios provide explicit mechanisms whereby a scale invariant spectrum of adiabatic perturbations may be generated without the need for cosmic inflation, albeit in a phase preceding the hot big bang singularity. A key question they face is whether there exists a unique prescription for following perturbations through the bounce, an issue which is not yet definitively settled. This goal of this paper is more modest, namely to study a bouncing Universe model in which neither General Relativity nor the Weak Energy Condition is violated. We show that a perturbation which is pure growing mode before the bounce does not match to a pure decaying mode perturbation after the bounce. Analytical estimates of when the comoving curvature perturbation varies around the bounce are given. It is found that in general it is necessary to evaluate the evolution of the perturbation through the bounce in detail rather than using matching conditions. ",
        "introduction": "In the inflationary Universe scenario, large scale structure is generated by the stretching of subatomic scale quantum fluctuations to macroscopic length scales, during an epoch of superluminal expansion supposed to have occurred prior to the radiation dominated era. Indeed such expansion appears to be essential to the creation of correlations on super-Hubble radius scales, a feat prohibited by causality in the standard hot big bang era.  However, cosmic inflation does not resolve the problem of the initial singularity, and it seems clear that a more complete theory than inflation is needed to deal with it. But if time existed before the the initial singularity, it is a logical possibility that the large scale structure and density perturbations we see today were generated during this pre-singularity epoch. The ekpyrotic and cyclic scenarios\\cite{ekpyrotic,cyclic} provide explicit realisations of this idea (pre-figured in the `pre-big bang' scenario of Veneziano et al.\\cite{pbb}). In the ekpyrotic and cyclic scenarios, scale invariant large scale density perturbations are generated during a phase which is contracting (in the Einstein frame) before the bounce to expansion\\cite{ekpyrotic}. The task of matching these perturbations across the bounce to the expanding epoch is highly challenging, since general relativity must break down there, and the use of string theory methods will ultimately be essential. There are indications that for the particular type of singularity involved here (namely the collapse of a single extra dimension) the divergences are relatively weak\\cite{ekperts}, but the matching issue remains unsettled at the present time\\cite{ekperts2,MartinPeter,Lyth}.  In a contracting universe, the growing mode adiabatic density perturbation corresponds to a shift in the time to the `big crunch'. A perturbation in the space curvature, however, is a decaying mode perturbation because it becomes increasingly irrelevant as the scale factor shrinks. (These statements both hold in Einstein frame and comoving gauge).  However, in an expanding universe, the situation is reversed. A time delay is now a decaying mode perturbation and the curvature perturbation is the growing mode. This has led several authors to suggest that in the ekpyrotic and cyclic scenarios, where scale invariant density perturbations are generated in the collapsing phase growing mode, that these perturbations would match on to pure decaying mode perturbations in the subsequent expanding phase. Indeed this result is obtained if one insists on matching the curvature perturbation across the bounce \\cite{ekperts2}.  However, this behaviour would appear surprising from a more physical viewpoint. In the models being considered, gravity is attractive throughout. The growing mode instability is driven by gravitational attraction and it is physically implausible that it should precisely reverse at a bounce so as to dissipate after the bounce and render the final universe perfectly homogeneous.  In this paper we address this issue by considering a model in which the issue can be definitively settled within conventional general relativity. It is well known that a massive scalar field in a closed Friedman-Robertson-Walker (FRW) Universe can lead to a bounce without a singularity, and without violating the Weak Energy Condition. This model has a long history starting with Schr\\\"{o}dinger \\cite{Schrodinger}, however he did not include the back-reaction of the scalar field on the background scale factor.  It was also employed in semi-classical studies of quantum effects (see for example \\cite{Fulling}), and of course is a common model in quantum cosmology and studies of the no boundary proposal (see for example \\cite{Hawking}, and more recently \\cite{Gratton}). The classical dynamics for the spatially homogeneous case have also been investigated quite extensively, see for example \\cite{bgrefs}.  A preliminary study of cosmological perturbations in such a Universe (and in other bouncing models) was performed by Hwang and Noh \\cite{Hwang}, with inconclusive results since the usual large scale approximations break down around the bounce. We develop new approximations that do not break down around the bounce, and using numerical simulations, accurately determine the propagation of linear perturbations through the bounce. We show that growing mode perturbations developed in the collapsing phase do not match to pure decaying mode perturbations in the expanding phase, thus there is no contradiction with the physical argument given above. Of course, this does not at all prove that the same is true in the ekpyrotic/cyclic models: as the original papers made clear, an unambiguous matching condition is still required. But the present work does show that propagation of a growing mode perturbation across a bounce is possible at least in this toy model.    We would also like to briefly mention recent work on the ekpyrotic/cyclic scenarios. Durrer and Vernizzi \\cite{Durrer} showed how matching across the bounce depends on what surface the matching is done on and the intrinsic stress energy of that surface. They showed that only for the special case of matching on constant energy density surfaces with no surface tension that the growing mode in the collapsing phase matches completely onto the decaying mode in the expanding phase. Generally, some of the growing mode in the collapsing phase will go to the growing mode in the expanding phase. We should also mention a recent study of the possibility of isocurvature perturbations in the Ekpyrotic model \\cite{Riotto}. Peter and Pinto-Neto \\cite{Peter1} studied the matching problem in a hydro-dynamical fluid, and in a second publication employed a scalar field with negative kinetic energy in order to obtain a regular bounce \\cite{Peter3}. The problem of matching quantum fields across a bouncing universe of the type relevant to the ekpyrotic/cyclic scenarios has been considered by Tolley and Turok \\cite{Tolley}, and a number of studies have been made of analogous bouncing models within string theory \\cite{stringmodels}. ",
        "conclusions": "In this article we have examined the way perturbations evolve across a bounce in a scalar field dominated closed Universe. The background model naturally bounces without violating General Relativity or the Weak Energy Condition. Analytical approximations were found for the background and perturbations. The perturbation approximations were found to work well around the bounce $|t|\\lesssim a_0$ and away from the bounce in the slow-roll regime $|t|\\gtrsim a_0$. However the overlap between the approximations was not sufficient to use the analytical approximations alone for arbitrary initial conditions. But they were useful in checking and interpreting the analytical solutions. The numerical solutions were complicated by the ordinary differential equation for $\\Phi$ having a singular coefficient at the bounce. Thus a Taylor approximation had to be used around the bounce and shooting methods used to obtain the numerical solution for particular initial conditions for a wider domain. The Taylor expansion around the bounce could also be used to show that the perturbations and their first and second time derivatives are finite at the bounce and so non-linearities in the Einstein tensor are not inevitable, but depend on the choice of initial conditions.  It was shown that all wave number modes violate the large scale condition for $|t|\\gtrsim a_0$ and so become oscillatory around the bounce. We also found during the collapsing phase the growing mode could cause the comoving pressure to increase sufficiently to make $\\R$ vary significantly as the bounce is approached. The same is true for the decaying mode as the bounce is approached from the expanding phase side.  Thus, in general we do not expect $\\R$ to be constant around the bounce. A formula predicting when $\\R$ starts to vary was given.  We also showed numerically that, for this model, a pure ingoing growing mode ($\\R^-=0$) did not lead to a pure outgoing decaying mode ($\\R^+=0$).  This model shows that in general the dynamics of the bounce need to be taken into account when evaluating how perturbations change across a bounce.   {\\bf Acknowledgements} We would like to thank M. Bucher, D. Lyth and D. Wands for helpful discussions. This research is funded by PPARC (UK)."
    },
    "0206/astro-ph0206269_arXiv.txt": {
        "abstract": "We developed a magneto-turbulent model for the cosmic ray (CR) electrons seen in the radio halo clusters of galaxies. Steady state momentum distribution functions of the CR electrons are calculated for given spectra of the turbulent Alfv\\'{e}n waves. The radio spectrum produced by the obtained CR electron distribution is compared to the observed radio spectrum of the Coma radio halo. We find that the observed radio spectrum of the Coma cluster is well reproduced when the spectral index of the  turbulent Alfv\\'{e}n waves is $\\sim 2.8$. The obtained energy spectrum of the turbulent Alfv\\'{e}n waves is steeper than that expected from the turbulence theory, suggesting back reaction of the particle acceleration. The fitting procedure constraints the amplitude of the turbulent Alfv\\'{e}n waves as well as the spectral index. Then we estimate the dissipation of the turbulent Alfv\\'{e}n waves, which is found to be less than the cooling rate by X-ray radiation. We suggest that the turbulence which is sufficient for particle acceleration is developed in the clusters containing the radio halo. It is most likely that cluster mergers create the turbulence and seed relativistic electrons. ",
        "introduction": "The radio halo in a cluster of galaxies is a diffuse nonthermal (synchrotron) radio emission seen in the central region of $\\sim 1$ Mpc in diameter. According to Hanisch (1982), the typical luminosity is $ L_{\\rm t} \\sim 10^{40} - 10^{41}$ erg s$^{-1}$ (10 MHz $\\leq \\nu \\leq$ 10 GHz), and the spectral index is $1 \\lesssim \\alpha \\lesssim 2$. The typical spatial scale of the radio halos is larger than the optical core radius, but smaller than the Abell radius. Giovannini, Tordi, \\& Feretti (1999) inspected a sample of 205 clusters from the X-ray-brightest Abell-type clusters (Ebeling et al. 1996) to search for new radio halo and relic candidates. They found only 29 candidates. In addition, they found that occurrence of the radio halos and relics is higher in clusters with higher X-ray luminosity and higher temperature and confirmed that the positive correlation between the absence of a cooling flow and the presence of radio halo. The rarity of the radio halo is one characteristic to be considered in its formation scenario.  In clusters of galaxies with radio halos, nonthermal X-ray radiation due to inverse Compton scattering of cosmic microwave background (CMB) photons is expected (Rephaeli 1979). Indeed, nonthermal hard X-ray radiation was recently detected in a few rich clusters (e.g., Fusco-Femiano et al. 1999; Rephaeli, Gruber, \\& Blanco 1999; Fusco-Femiano et al. 2000) and several galaxy groups (Fukazawa 2001), although their origin is still unidentified.  Among clusters with radio halo, the Coma cluster is the most informative. The halo size is $30'-40'$ in diameter corresponding to $600-800\\ h_{100}^{-1}$ kpc. The radio halo region contains two bright radio galaxies, NGC4874 and NGC4869. The spatially averaged spectral index of the halo, with reduction of contamination by the radio galaxies, is $\\alpha \\sim 1.3$. Giovannini et al. (1993) presented a spatial distribution of the spectral index between 326 MHz and 1.4 GHz. They found that the radio spectrum of the central halo region with size of $\\sim 15'$ is flatter, $\\alpha \\sim 0.8$, while in the vicinity of the tailed radio source associated with NGC4869 is steeper, $\\alpha \\sim 1.8$. This means that high energy cosmic rays (CRs) are more abundant in the halo region than in the vicinity of the tailed source. The flatness of the radio halo spectrum suggests particle acceleration in the intracluster space. This has been pointed out by Giovannini et al. (1993). The radio spectrum of the outer region of the radio halo is steeper ($\\alpha \\sim 1.8$) than that of the central region.  Kim et al. (1990) studied the magnetic fields in the Coma cluster by using the rotation measures (RMs) of background radio sources (QSO and radio galaxies). The observed RMs of the background sources seen through the Coma cluster have an excess of $\\sim$ 38 rad m$^{-2}$. By adopting a galaxy scale of $10-40$ kpc as a correlation length, they found the (electron density weighted) amplitude of the random magnetic field to be $\\sim$ 1 $\\mu$G. On the other hand, the comparison between nonthermal hard X-ray emission and synchrotron radio emission gives the volume averaged magnetic field strength, if all the hard X-ray is emitted through inverse Compton scattering of CMB photons (Rephaeli 1979). Rephaeli, Gruber, \\& Rothschild (1987) obtained a lower limit value $\\sim 0.11~\\mu$G by using the upper limit on the hard X-ray emission from the Coma cluster. The first detection of the hard X-ray emission from the Coma cluster by {\\it BeppoSAX} gives the field strength of 0.15 $\\mu$G (Fusco-Femiano et al. 1999). The obtained values are significantly smaller than those derived from RMs. However, it is important to note that (1) these estimates assume that the emitting volumes of the hard X-ray and radio emission coincide, which might not be guaranteed, and that (2) these estimates are affected by the presence of additional infrared background radiation fields (Schlickeiser \\& Rephaeli 1990). These simplification can make the derived field strength smaller than that derived from RMs.  The origin of the CR electrons in the clusters with radio halo is still unclear. Jaffe (1977) proposed the primary electron model where the radio halo is an emission by CR electrons diffusing away from radio galaxies in a cluster. However, this model contains some difficulties, one of which is that the diffusion distance does not seem as large as the radio halo extent. Dennison (1980) proposed the secondary electron model where the relativistic electrons are produced through decay of charged pions induced by the interaction between relativistic protons from radio galaxies and thermal protons in ICM. In this model, however, too much gamma-ray emission is produced at least for the Coma cluster (Blasi \\& Colafrancesco 1999).  Excess of the high energy CR electrons in the halo region, indicated by the flat spectrum, strongly suggests the particle acceleration in an intracluster space. Models considering particle acceleration by intracluster magnetic turbulence were discussed by Roland (1981), Schlickeiser, Sievers, \\& Thiemann (1987), and Petrosian (2001). In this paper, we model the radio halo in terms of particle acceleration by the intracluster turbulent magnetic fields. The turbulence is assumed to be an ensemble of Alfv\\'{e}n waves, and the CR electrons are accelerated by pitch angle scattering by the Alfv\\'{e}n waves. Distribution functions of the CR electrons in energy space are obtained by solving a Fokker-Planck equation for the assumed turbulent spectra with various power indices. The calculated radio spectra are compared with the observed one to determine the energy spectrum of the turbulent Alfv\\'{e}n waves.  Deiss et al. (1997) proposed a merger shock acceleration model which explains the rarity of radio halos. Numerical simulations by Takizawa \\& Naito (2000) show radio emission is luminous while merger shocks accelerate CR electrons but decrease rapidly because of inverse Compton cooling after the shocks disappear. The transient feature of this model is consistent with the rarity of radio halos. We will discuss the relation between our model and theirs.  In \\S2, the basic equations to calculate a radio spectrum are introduced. In \\S3, the calculated spectra are interpreted, and the energy spectra of the turbulent Alfv\\'{e}n waves is determined so as to reproduce the observed radio spectrum. In \\S4, the generation mechanism of the obtained energy spectra is investigated, and over all turbulence structure is discussed. \\S5 is for summary. ",
        "conclusions": "\\subsection{The Energy Spectrum of Turbulent Alfv\\'{e}n Waves} Let us consider physical status of the turbulent Alfv\\'{e}n waves in our best fit model. Ruzmaikin \\& Shukurov (1982) gives the rate of energy-transfer via the nonlinear interaction among different wave numbers as \\begin{eqnarray} \\left ( \\frac{dE}{dt} \\right )_{\\rm nl} & \\sim & \\left. \\frac{P(k)^2 k^3}{\\rho v_{\\rm A}}, \\right.\\label{wavenl} \\end{eqnarray} where $\\rho$ is the mean gas density. If the energy spectrum of the turbulent Alfv\\'{e}n waves is determined only through this energy cascade process, it becomes $P(k) \\propto k^{-3/2}$ in a steady state because $(dE/dt)_{\\rm nl}$ is constant for $k$ (Ruzmaikin \\& Shukurov 1982). However, this is significantly flatter than that in our best fit model whose exponent is $-2.8$. Thus, other physical processes might work to steepen the Alfv\\'{e}n wave spectrum. Obviously, the Alfv\\'{e}n waves lose their energy through the particle acceleration. The CR electrons between $p$ and $p+dp$ absorb the energy of Alfv\\'{e}n waves at the rate, \\begin{eqnarray} \\left ( \\frac{dE}{dt} \\right )_{\\rm accel} &   =  & 4 \\pi p^2 (m_{\\rm e0} c^2) \\dot{\\gamma}_{\\rm accel} f(p) dp  \\nonumber \\\\ & \\sim & 4 \\pi (m_{\\rm e0} c^2) p^3 \\frac{\\gamma}{T_{\\rm a}}f(p) \\nonumber \\\\ & \\sim & 4 \\pi m_{\\rm e0}^4 c^5 (w+2) a_1 \\gamma^{w+2} f(p), \\label{waveacc} \\end{eqnarray} where we use $dp \\sim p$, $\\dot{\\gamma}_{\\rm accel}=\\gamma/T_{\\rm a}$. We checked whether this plays a significant role or not in our best fit model. For simplicity, we approximate $P(k) \\propto k^{-3}$ and $f(p) \\propto p^{-6}$. In this case, $(dE/dt)_{\\rm nl} \\propto k^{-3}$ and $(dE/dt)_{\\rm accel} \\propto k^{-1}$ because the resonant wave number is proportional to $p^{-1}$. Therefore, the energy loss due to the particle acceleration is more significant in higher wave numbers and may be attributed to the spectral steepening. We checked if both terms are comparable to each other in the observed frequency range (30 MHz $-$ 5 GHz), which corresponds to the wave number range of ($10^{-14} \\sim  10^{-13}) B_{\\mu {\\rm G}}^{3/2}$ cm$^{-1}$ (eq.[\\ref{k-g}]). The critical wave numbers $k_{\\rm t}$ where $(dE/dt)_{\\rm accel} = (dE/dt)_{\\rm nl}$ are shown in Table \\ref{tab3} for our best fit models. They are certainly within or near the wave number range mentioned above. The obtained steep wave spectra, $w \\sim$ 2.8 and 4.5 for the Coma cluster and $w > 2.0$ in general, could be brought about by this effect. We suggest that time dependent and self-consistent calculations for non-linear wave-particle systems with back reaction of particle acceleration are necessary to investigate evolution of the systems.  Next, we make a crude estimate of $P(k)$ in $k<k_{\\rm t}$ and check the consistency of our consideration mentioned above. We adopt $P(k)$ of our best fit models for $k>k_{\\rm t}$ and assume $P(k) \\propto k^{-\\xi}$ for $k<k_{\\rm t}$. As a result, the broader band model power spectrum is, \\begin{eqnarray} P(k) = \\left\\{ \\begin{array}{@{\\,}ll} P_{\\rm t} \\biggl( \\frac{k}{k_{\\rm t}} \\biggr)^{-\\xi} &  \\mbox{($k_{\\rm min} < k < k_{\\rm t}$)} \\\\ P_{\\rm t} \\biggl( \\frac{k}{k_{\\rm t}} \\biggr)^{-w} &  \\mbox{($k_{\\rm t} < k$)} \\end{array} \\right. , \\end{eqnarray} where $P_{\\rm t} \\equiv b_{\\rm t} (B_0^2/8 \\pi)(w-1)/k_{\\rm t}$ and $k_{\\rm min}$ is the wave number corresponding to the scale of the cluster, $k_{\\rm min} \\sim 10^{-24}$ cm$^{-1}$. Then, the total energy density of Alfv\\'{e}n waves $M_{\\rm tot}$ is given by \\begin{eqnarray} M_{\\rm tot} =  \\frac{w-1}{\\xi-1} b_{\\rm t}  \\frac{B_0^2}{8 \\pi} \\biggl(\\frac{k_{\\rm t}}{k_{\\rm min}} \\biggr)^{\\xi-1} \\biggl\\{ 1 - \\frac{w-\\xi}{w-1} \\biggl( \\frac{k_{\\rm t}}{k_{\\rm min}} \\biggr)^{1-\\xi} \\biggr\\}, \\end{eqnarray} where we assumed $\\xi > 1$. If we assume $M_{\\rm tot}$ is as large as the energy density of the large-scale-field $B_0^2/(8\\pi)$, then we obtain $\\xi$ with given $k_{\\rm t}$ and $b_{\\rm t}$. When $B_0=1\\ \\mu$G, $k_{\\rm t}\\sim10^{-14}$ cm$^{-1}$ and $b_{\\rm t} \\sim 10^{-6}$, we have $\\xi \\sim 1.6$, which is close to the index of the energy spectrum, $3/2$, formed by the non-linear interaction. Therefore, this result is consistent with the previous picture of the non-linear interaction for $k < k_{\\rm t}$ and the acceleration reaction for $k > k_{\\rm t}$. When $B_0=0.1\\ \\mu$G, $k_{\\rm t}\\sim10^{-15}$ cm$^{-1}$ and $b_{\\rm t} \\sim 10^{-3}$, we have $\\xi \\sim 1.2$.  The rate of the energy transfer from $k_{\\rm 30MHz}$ to higher wave numbers is estimated as $\\varepsilon \\sim [P(k_{\\rm 30MHz})]^2 k_{\\rm 30MHz}^3 / (\\rho v_{\\rm A})$. Table \\ref{tab3} shows that the values of $\\varepsilon$ are less than the thermal X-ray cooling rate except for the unrealistic case with $B_0=0.1~\\mu$G, $w=4.5$ and $\\nu_{\\rm c} = 10$MHz. Therefore, even if the energy of the turbulent Alfv\\'{e}n waves goes into the thermal energy of the hot gas by the dissipation at higher wave numbers, the hot gas could not be heated up.    \\subsection{The Origin of the Turbulence} The energy transfer rate $\\varepsilon_0$ at $k_{\\rm min} \\sim 10^{-24}$ cm$^{-1}$ is estimated as $\\varepsilon_0 \\sim P_0^2 k_{\\rm min}^3 /\\rho v_{\\rm A} \\sim 10^{-31} B_{\\mu {\\rm G}}^3$ [erg cm$^{-3}$ s$^{-1}$], where $P_0 \\sim B_0^2/(8 \\pi k_{\\rm min})$. Thus, an energy input rate higher than $ \\varepsilon_0$ is necessary to excite and maintain the turbulence. It is certain that motion of galaxies through the intracluster medium can excite turbulent eddies. In this case, however, the energy input rate of the steady turbulence is \\[ \\epsilon_{\\rm input} \\sim \\rho v_0^3 / l_0 \\sim 10^{-33} \\  \\mbox{erg} \\mbox{ s}^{-1} \\mbox{ cm}^{-3}, \\] where $l_0 \\sim 10$ kpc is a dimension of the galaxy, and $v_0 \\sim 3.2$ km s$^{-1}$ is the turbulent velocity at $l_0$ (Goldman \\& Rephaeli 1991). This input rate is insufficient for the required $\\varepsilon_0$. If the energy of the turbulence is input by the system of scale $R \\sim 1$ Mpc with the velocity comparable to the sound velocity $v_{\\rm s} \\sim 1000$ km s$^{-1}$, the energy input rate is \\begin{eqnarray} \\epsilon_{\\rm input} &\\sim& \\frac{\\rho v_{\\rm s}^2} {R/v_{\\rm s}} \\nonumber \\\\ &\\sim& 10^{-28} \\left( \\frac{n}{10^{-3} \\mbox{cm}^{-3}}\\right) \\left( \\frac{v_{\\rm s}} {1000\\mbox{ km s}^{-1}} \\right)^3 \\nonumber\\\\ &    & \\times \\left( \\frac{R}{1\\mbox{ Mpc}} \\right)^{-1} \\mbox{ erg} \\mbox{ s}^{-1} \\mbox{ cm}^{-3}. \\end{eqnarray} This input rate is sufficient for the radio halo formation. A merging event between sub-structures can input the kinetic energy into the turbulence at this high rate. Moreover, cluster merger can produce seed CR electrons through 1st order Fermi acceleration at shocks in ICM (Takizawa \\& Naito 2000). Indeed, recent high resolution MHD simulations of cluster mergers (Roettiger et al. 1999) show that the bulk flow is replaced by turbulent motion in the later stages of the merger.  \\subsection{The Supply of Cosmic Ray Electrons} We calculate the total number of the CR electrons $N_{\\rm CR}$ from the luminosity and find $N_{\\rm CR} \\sim 10^{60}$ for the Coma radio halo, where we adopt the luminosity $L_{\\rm t} \\sim 10^{41}$ erg s$^{-1}$ (10 MHz $ \\leq \\nu \\leq $ 10 GHz), the spectral index $\\alpha \\sim 1.3$, and the magnetic field strength $B_0 \\sim 1~\\mu$G.  We examine how much the radio galaxy can supply the CR electrons. NGC4874 is one of the two dominant galaxies in the Coma cluster. The observed flux of the radio source associated with NGC4874 is $S_{\\rm 408MHz} \\sim 4 \\times 10^{-24}$ erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$ at 408 MHz (Jaffe \\& Perola 1974). From the observed flux  we calculate the number density of the CR electrons at the injection energy $\\sim 100$ MeV. For simplicity, we assume that the shape of the source is a sphere with a radius $r \\sim 10$ kpc, and that the radio spectrum follows a single power law with an index $\\alpha \\sim 0.5$. We also assume the magnetic field strength is 10 $\\mu$G. We find the number density $n_{\\rm CR} \\sim 6 \\times 10^{-11}$ cm $^{-3}$ at 100 MeV. The injection rate of the CR electrons leaked from the radio source is given by $\\dot{n}_{\\rm CR} \\sim 4 \\pi r^2 n_{\\rm CR} v_{\\rm D}$, where $v_{\\rm D}$ is the diffusion velocity of the CR electrons. Taking $v_{\\rm D} \\sim v_{\\rm A} \\sim 700$ km s$^{-1}$ gives an upper limit of the injection rate $\\dot{n}_{\\rm CR} \\lesssim 4 \\times 10^{43}$ s $^{-1}$. When  $\\dot{n}_{\\rm CR} \\sim 4 \\times 10^{43}$ s $^{-1}$, the CR electrons in the radio halo are supplied in $N_{\\rm CR}/\\dot{n}_{\\rm CR} \\sim 7 \\times 10^8$ yr.  The merger shock acceleration supplies the CR electrons in the radio halo region. Takizawa \\& Naito (2000) assumed that the total kinetic energy of accelerated electrons is 5 \\% of the viscous energy, which is nearly equal to the energy of the shock heating, and they found that the luminosity of the synchrotron radio emission (10 MHz$-$10 GHz) is $\\sim 10^{43}$ erg s$^{-1}$. We find the number of the CR electrons calculated from the luminosity is $5 \\times 10^{62}$ and sufficient for the observed radio halo, where we assume the magnetic field strength is $0.1~\\mu$G and the radio spectral index is 0.7. According to Takizawa \\& Naito (2000), the radio luminosity decreases after the most contracting epoch. We suggest that the turbulent Alfv\\'{e}n waves reaccelerate the CR electrons injected by the merger shock and maintain the radio halo."
    },
    "0206/astro-ph0206096_arXiv.txt": {
        "abstract": "We present theoretical \\ion{Fe}{2} emission line strengths for physical conditions typical of Active Galactic Nuclei with Broad-Line Regions. The \\ion{Fe}{2} line strengths were computed with a precise treatment of radiative transfer using extensive and accurate atomic data from the Iron Project.  Excitation mechanisms for the \\ion{Fe}{2} emission included continuum fluorescence, collisional excitation, self-fluorescence amoung the \\ion{Fe}{2} transitions, and fluorescent excitation by \\lya\\ and \\lyb. A large \\ion{Fe}{2} atomic model consisting of 827 fine structure levels (including states to $\\rm E\\,\\approx 15\\;eV$) was used to predict fluxes for approximately 23,000 \\ion{Fe}{2} transitions, covering most of the UV, optical, and IR wavelengths of astrophysical interest.  Spectral synthesis for wavelengths from $\\lambda\\,1600\\,$\\AA\\ to $1.2\\,\\mum$ is presented. Applications of present theoretical templates to the analysis of observations are described. In particular, we discuss recent observations of near-IR \\ion{Fe}{2} lines in the $8500\\,$\\AA--$1\\,$\\mum\\ region which are predicted by the \\lya\\ fluorescence mechanism. We also compare our UV spectral synthesis with an empirical iron template for the prototypical, narrow-line Seyfert galaxy I~Zw~1.  The theoretical \\ion{Fe}{2} template presented in this work should also applicable to a variety of objects with \\ion{Fe}{2} spectra formed under similar excitation conditions, such as supernovae and symbiotic stars. ",
        "introduction": "Transitions of singly ionized iron dominate the spectra of many astrophysical objects, from the sun and stars to active galactic nuclei (AGNs) and quasars (Viotti 1988). However, the interpretation of this spectrum is complex, and to extract meaningful results for the physical conditions in the emitting region, and the iron abundance and ionization fractions, one is faced with the solution of a complex radiative transfer problem requiring the specification of many thousands of radiative and collisional rates in a non-local thermodynamic equilibrium (non-LTE) formalism.  Until recently, such calculations have been hampered by the paucity of basic atomic data for \\ion{Fe}{2}.  However, the Iron Project (Hummer \\etal\\ 1993) has been specifically initiated to address this problem, and new, accurate atomic data for \\ion{Fe}{2} have been calculated. In particular, radiative dipole transition probabilities for over 21,000 fine-structure transitions of \\ion{Fe}{2} have been computed by Nahar (1995), and collision strengths for over 11,000 fine-structure transitions have been computed by Zhang \\& Pradhan (1995) and Bautista \\& Pradhan (1996). These calculations, and the Iron Project in general, employ the powerful and accurate \\underline{R}-matrix method (Burke \\& Berrington 1993).  In our earlier work on \\ion{Fe}{2} (Sigut \\& Pradhan 1998; SP98), we employed a limited, non-LTE atomic model with 262 fine structure levels which was still sufficiently large for \\lya\\ fluorescent excitation to be investigated in detail. It was shown that \\lya\\ excitation can be of fundamental importance in enhancing the UV and optical \\ion{Fe}{2} fluxes. In particular, it was predicted that \\lya\\ fluorescence results in significant near-infrared \\ion{Fe}{2} emission in the region $\\lambda\\lambda\\,8500-9500\\,$\\AA. Following the SP98 work, recent observations have detected many of these near-IR \\ion{Fe}{2} emission lines from several narrow-line Seyfert~I galaxies (Rodriguez-Ardila \\etal 2001), and from a Type~IIn supernova remnant with narrow emission lines (Fransson \\etal 2001).  Although difficult to observe, these near-IR \\ion{Fe}{2} lines should be indicative of the excitation mechanisms and the possible interplay between collisional and fluorescent excitation (Rudy \\etal 2000).  As our predicted near-IR \\ion{Fe}{2} fluxes in this wavelength region are likely to be of wider interest, this paper presents a detailed line list from our non-LTE calculations with an extended \\ion{Fe}{2} model atom.  We intend this line list to be our first step in developing a reliable set of {\\it theoretical\\/} templates for the iron emission from AGN. Currently, due to the complexity of the observed iron emission from AGN, such emission is typically modeled using empirical templates derived from specific AGN spectra (Boroson \\& Green 1992, Corbin \\& Boroson 1996). A recent example of this method is the \\ion{Fe}{2}-\\ion{Fe}{3} template of Vestergaard \\& Wilkes (2001) derived from high-quality UV spectra of the narrow-line Seyfert~I galaxy I~Zw~1.  Such templates play a critical role in extracting a measure of the total iron emission from heavily blended and broadened AGN spectra.  For example, Dietrich \\etal (2002) apply the Vestergaard \\& Wilkes template to extract a measure the relative iron-to-magnesium abundance ratio from a sample of high-$z$ quasars. Such studies seek to constrain the epoch of major star formation in AGN using the iron-to-magnesium abundance ratio as a nucleosynthesis ``clock\" following Hamann \\& Ferland (1992). Such studies are beginning to impose important cosmological constrains: for example, Aoki, Murayama \\& Denda (2002) have detected \\ion{Fe}{2} emission from a $z=5.74$ QSO with a strength comparable to much lower redshift objects.  Empirical templates have the advantage that they can side-step the complicated process of specifying in detail the iron emission mechanisms, and have generally found to provide better fits to observations than theoretical templates (Iwamuro \\etal\\ 2002, Thompson, Hill \\& Elston 1999).  Nevertheless, there is still a strong need to develop reliable theoretical templates:  (1) empirical templates assume that the underlying AGN population used to construct the template is typical and that the iron emission in other related objects can be modeled as a simple scaling of the fiducial spectrum.  (2) Empirical templates can never be completely free of the complications introduced by the large blending and broadening present in AGN spectra. For example, it is difficult to constrain the iron emission present at the location of the \\ion{Mg}{2} h \\& k lines, although such an estimate does affect the derived fluxes.  Theoretical iron flux templates can address both of these problems, allowing estimates of the response of the iron emission to model parameters which many differ from object to object (such as the photoionizing radiation field), and providing a spectrum from which complex blended features can be decomposed into their individual contributions. But theoretical templates must first show that they can explain the current empirical templates.  In this work, we will compare our predictions with the empirical UV template of Vestergaard \\& Wilkes.  \\subsection{The Physics of \\protect\\ion{Fe}{2} Line Formation in AGN}  The formation of the \\ion{Fe}{2} emission spectrum from AGN is still poorly understood (Joly 1993, Hamann \\& Ferland 1999).  Typically, photionization cloud models for the BLR fail to account for the observed strength of the \\ion{Fe}{2} emission. A class of ``super-strong'' \\ion{Fe}{2} emitters is known (Lipari, Macchetto \\& Golombek 1991; Graham \\etal\\ 1996) which seem to be unaccountable by traditional photoionized models.  In these cases, and possibly all, a different cloud population may be the origin of the \\ion{Fe}{2} emission, such as mechanically heated clouds shielded from the central continuum source (Joly 1987), perhaps originating in the outer regions of an accretion disk (Collin-Souffrin \\etal\\ 1988).  An important sub-class of AGN with BLRs are the narrow-line Seyfert galaxies (Osterbrock \\& Pogge 1985). I~Zw~1 is the prototypical narrow-line quasar (Laor \\etal 1997), and it is also a strong \\ion{Fe}{2} emitter (Marziani \\etal 1996). Such narrow-line Seyfert~1 galaxies (NLS1) enable both better emission line diagnostics and better tests of theoretical spectra as their spectra are broadened with typical velocities of $\\leq 1000\\,\\rm km\\,s^{-1}$.  The accepted micro-physics of \\ion{Fe}{2} line formation in AGN is that of Wills, Netzer \\& Wills (1985) with the extension by Elitzur \\& Netzer (1985) to include fluorescent excitation by \\lya. The \\lya\\ excitation process is further studied by Johansson \\& Jordan (1984), Penston (1987), Sigut \\& Pradhan (1998), and Verner \\etal (1999).  The proposed excitation mechanisms can be understood with the aid of the highly simplified \\ion{Fe}{2} energy level diagram shown in Fig.~\\ref{fig:fe2simp}, based on a similar figure by Penston (1987). Four principal excitation mechanisms have been included:  \\vspace{3mm}  (1) {\\sc Continuum Fluorescence}: Photons incident on the illuminated face of the BLR cloud are absorbed in the resonance transitions and are subsequently re-radiated in the resonance and optical lines. Strong optical emission, however, requires thermalization of the resonance transitions in order to shift the effective branching ratio towards optical emission. Thus this mechanism suffers from having the photon source in the wrong location, namely outside the cloud at small optical depth, something first noted by Netzer (1988).  (2) {\\sc Collisional Excitation}: Inelastic collisions with electrons excite the odd parity levels near 5~eV which then decay into the optical and UV lines.  This mechanism is efficient whenever the gas temperature is above $\\approx 7000\\,$K, temperatures which are generally found in photoionized models of the BLR. Excitation is irrespective of the local optical depth in the \\ion{Fe}{2} lines, and thus this mechanism does not suffer the limitations of continuum fluorescence.  It is generally believed that collisional excitation is responsible for the bulk of the \\ion{Fe}{2} emission.  (3) {\\sc Self-Fluorescence}:  Netzer \\& Wills (1983) suggested that self-fluorescence, that is absorption of the \\ion{Fe}{2} UV resonance photons by overlapping UV \\ion{Fe}{2} transitions originating from the odd parity levels near 5~eV (labeled ``unexpected UV\" in Fig.~\\ref{fig:fe2simp}) was an important source of excitation to highly excited states due to the large number of wavelength coincidences between these groupings of levels.  (4) {\\sc Fluorescent excitation by} \\lya:  Penston (1987) noted that, despite theoretical calculations to the contrary (Elitzur \\& Netzer 1985), there is indirect evidence that \\lya\\ fluorescence may be an important but overlooked excitation mechanism. Penston noted the presence of unexpected UV \\ion{Fe}{2} lines (see Fig.~\\ref{fig:fe2simp}) in the spectrum of the symbiotic star RR Tel that seemed attributable only to cascades from higher levels pumped by \\lya\\ fluorescence. The emission nebulae of symbiotics offer densities and ionization parameters similar to those inferred for the BLRs of AGN.  Graham \\etal\\ (1996) have identified emission from the UV \\ion{Fe}{2} multiplets expected to be preferentially strengthened by this mechanism, as noted by Penston, in the spectrum of the ultra-strong \\ion{Fe}{2} emitter 2226-3905.  \\vspace{3mm}  In this current work, in addition to considerably enlarging the \\ion{Fe}{2} atomic model to 827 fine structure levels up to $\\rm E\\approx 15\\,$eV, we make improvements to the modeling of all of these excitation mechanisms:  (1) We solve the equation of radiative transfer with continuum fluorescence included through the appropriate boundary conditions on the transfer equations.  (2) We use a large and accurate set of {\\underline R}-matrix collision strengths for electron impact excitation of \\ion{Fe}{2}. Such rates are available for many of the key odd parity levels near 5~eV with are the upper levels of most of the UV and optical emission.  (3) Line-overlap amoung the \\ion{Fe}{2} transitions is included exactly in the radiative transfer solutions. We considerably expand the \\ion{Fe}{2} atomic model to 827 fine-structure levels and over 23,000 radiative transitions.  (4) We include frequency-dependent source functions for \\lya\\ and \\lyb\\ in the monochromatic source functions used in the radiative transfer solutions.  While the \\lya\\ and \\lyb\\ source functions used are approximate, as discussed in the Section~\\ref{sec:calclya}, their inclusion into the radiative transfer solution is exact. ",
        "conclusions": "Multi-level, accelerated lambda-operator techniques for non-LTE radiative transfer now allow solutions to be obtained for highly realistic atomic models including complex cases of line-overlap and fluorescent excitation. Coupled with the new atomic data from the Iron Project, such techniques have been applied for the first time with a reasonably complete \\ion{Fe}{2} atomic model to theoretical AGN BLR spectra. The theoretical \\ion{Fe}{2} line fluxes presented should help in the identification of \\ion{Fe}{2} transitions in AGNs and related sources, and in the delineation of excitation mechanisms producing the \\ion{Fe}{2} spectrum.  We are extending the calculations to include the line spectra of other iron ionization stages, principally Fe\\,{\\sc i} and Fe\\,{\\sc iii}. Laor \\etal\\ (1997) and Vestergaard \\& Wilkes (2001) specifically note the present of significant \\ion{Fe}{3} in the spectrum of I~Zw~1, and Graham \\etal\\ (1996) have detected Fe\\,{\\sc iii} emission in the ultra-strong Fe\\,{\\sc ii} emitter 2226-3905.  Kwan \\etal\\ (1995) have detected Fe\\,{\\sc i} emission in two Fe\\,{\\sc ii}-strong quasars, IRAS 07598+6508 and PHL 1092.  Simultaneous modeling of these ionization stages should provide more constraints on the nature of the iron emission.  Our calculations currently include transitions between observed energy levels whereas there remains a large number of theoretically predicted energy levels.  We are including these levels and the implied radiative transitions in order to provide a much more complete description of the Fe\\,{\\sc ii} emission spectrum.  We are currently working on bringing the entire photoionization calculation for all atoms and ions within the framework of exact radiative transfer established in this work.  This will allow a self-consistent treatment of the Fe\\,{\\sc ii} emission by including it in the net heating/cooling which determines the temperature structure, and will also allow an {\\em exact\\/} treatment of \\lya\\ fluorescent excitation of Fe\\,{\\sc ii} emission.  We are also working on several other specific problems, such as the interpretation of Fe\\,{\\sc ii}/Mg\\,{\\sc ii} line ratios, employing extended non-LTE models for the relevant atomic species and exact radiative transfer.  The tabular and graphical material presented in this work is available electronically on request.  \\vspace{0.3in}  \\noindent We would like to thank Sultana Nahar for numerous contributions, and Marianne Vestergaard for the data in Figure~\\ref{fig:1zw1_uv}.  This work was supported by the Natural Sciences and Engineering Research Council of Canada (TAAS), and by the U.S. National Science Foundation and NASA (AKP).   \\clearpage"
    },
    "0206/astro-ph0206119_arXiv.txt": {
        "abstract": "{ We observed the 3335 MHz ($\\lambda$ 9cm) F=1-1 line of CH toward a sample of diffuse clouds occulting compact extragalactic mm-wave continuum sources, using the old NRAO 43m telescope.  Because radiofrequency observations of CH really must be calibrated with reference to a known CH abundance, we begin by deriving the relationships between CH, \\EBV, \\HH\\ and other hydrides found by optical spectroscopy. No simple relationship exists between N(CH) and \\EBV, since N(CH) is strongly bimodal with respect to reddening for \\EBV $<$ 0.3 mag and the typical range in the N(CH)/\\EBV\\ ratio is an order of magnitude or more at any given \\EBV $> 0.3$ mag.  However, N(CH)/N(\\HH)  $=   4.3 \\pm 1.9 ~\\times~10^{-8}$ in the mean and N(CH) $\\propto$ N(\\HH)$^ 1.00\\pm0.06$ for $10^{19} < $ N(\\HH) $ < 10^{21}\\pcc $.  If CH is a good predictor of \\HH, 40\\%-45\\% of the hydrogen in the local diffuse/translucent ISM is in the molecular form at the accepted mean density, higher than previous estimates found in samples of lower-than-average mean density. Optical observations of the population ratios in the upper and lower halves of the CH lambda-doublet suggest that the brightness of the 3335 MHz CH line should be double-valued at a given CH column density in diffuse gas:  double-valuedness is noticeable in our data when comparing CH with CO or \\hcop. The CH brightness at 3335 MHz is mildly bimodal with respect to CO emission in our diffuse cloud data but much more strongly bimodal when comparing diffuse or translucent gas and dark gas.  The CH $\\Lambda$-doublet is generally inverted in diffuse gas but we did not succeed in measuring the excitation temperature except toward 3C123 where we confirm one older value $\\Texc \\approx -10$ K. ",
        "introduction": " ",
        "conclusions": "Although fourth in the current series, this paper is actually also the third of three papers discussing singledish cm-wave spectra of the molecules in diffuse gas, taken toward and around a sample of compact, extragalactic mm-wave continuum sources.  In H I, such on-off comparison experiments have come to be known as emission-absorption experiments \\citep{DicTer+78} but as it turns out, we performed one emission-absorption experiment (in OH; \\cite{LisLuc96}), one absorption-absorption experiment (in \\HH CO, which appeared in absorption both on and off-source; \\cite{LisLuc95a}) and, here, one emission-emission experiment (since no true absorption was detectable).  We began by displaying the richly structured behaviour of CH with \\EBV; N(CH) is multi-valued with respect to \\EBV, depending on the degree of conversion to molecular gas along the line of sight, and a simple, linear CH-\\EBV\\ relationship can be expected only when the extinction is dominated by molecular gas, as toward a single dark cloud.  Otherwise, the range of measured N(CH) at a given \\EBV\\ in the diffuse gas as a whole is typically more than one order of magnitude.  Much of this behaviour can be explained on the basis of an easily-demonstrated and long-known, nearly constant relative abundance $<$\\XCH$>$ = $4.3\\pm1.9\\times10^{-8}$, and N(CH) $\\propto$ N(\\HH)$^{1.00\\pm0.06}$ for N(\\HH) $\\la 10^{21}~\\pcc$: as well, we have that N(\\HH) $\\propto$ \\EBV$^{1.8}$ for 0.2 $<$ \\EBV\\ $<$0.7.  If CH is a good predictor of \\HH, the 140 lines of sight gathered to study the CH-\\EBV\\ relationship allow derivation of the molecular fraction in the diffuse/translucent ISM over a much wider range of sample mean densities $<$\\EBV$>$/$<$R$>$ than is directly accessible in measurements of the lines of hydrogen.  The molecular fraction found in this way is in good agreement with direct measurements at low (Copernicus) and high (FUSE) sample mean density, and is 0.4-0.45 for $<$\\EBV$>$/$<$R$> =  0.61$ mag kpc$^{-1}$, which is the accepted mean in the gas within 500-1000 pc.  We pointed out that sensitive optical measurements of the population ratio in the upper and lower halves of the ground-state CH $\\lambda$-doublet toward two stars predict that the brightness of the microwave CH lines should be double-valued at a given CH column density in diffuse gas depending on whether the excitation is inverted (the brighter branch); this is consistent with models of CH excitation which predict a transition from normal excitation to inversion at hydrogen densites in the range 10 - 1000 $\\pcc$, but the effect is not present in the microwave lines in these directions.  This could be due to the disparity in beam-sizes or to relatively small errors in the optical data.  We presented 3335 MHz CH observations toward some of the compact extragalactic mm-wave continuum sources studied in this series of papers, toward two strong cm-wave sources, and around \\zoph, and compared the properties of CH with those of OH and CO in emission and \\hcop\\ and \\cch\\ seen in absorption. In stronger-lined gas, the CH/OH comparison confirms the very small OH excitation temperatures which have been found in diffuse gas.  Comparisons of CH with \\hcop\\ and \\cch\\ show that there is either a very large scatter in the CH brightness or microwave-derived CH column density at a given N(\\hcop) or N(\\cch) or perhaps a bimodality.  The CH/\\hcop\\ comparison readily (but only roughly) confirms our previously-derived ratio N(\\hcop)/N(\\HH) $= 2 \\times 10^{-9}$.  The 3335 MHz line brightness in diffuse gas is very definitely bimodal with regard to CO emission in both the new data presented here and our previously-published data around \\zoph.  To explore this further, we compared our data with other published studies which used CH to derive the CO-\\HH\\ conversion factor in diffuse/translucent gas and found that they are entirely consistent with our data.  This may be a manifestation of a disparity between inverted and non-inverted CH expected in diffuse gas.  We found some hitherto-unnoticed systematic behaviour in the CH-CO comparison in diffuse and dark gas, in particular a steady, factor of $\\approx 3$ offset in the ratio of CH and CO profile integrals for W(CO) = 1 - 30 K \\kms:  W(CH)/W(CO) is consistently larger by $\\approx 3$ in dark gas. A shallow slope in the W(CH)-W(CO) relationship in diffuse gas is undertandable because the CO abundance varies rapidly with N(\\HH), N(CO) $\\propto $ N(\\HH)$^2$ and the CO brightness will increase even faster than N(CO), but the presence of nearly the same shallow slope W(CH) $\\propto W({\\rm CO})^{0.3}$ in dark and diffuse gas is puzzling. It may reflect the decline of \\XCH\\ which is known to occur in very dark gas.  The next paper in this series will discuss several species whose abundances are best determined at cm-wave frequencies, such as \\ammon, \\hhco\\ and C$_4$H."
    },
    "0206/astro-ph0206082_arXiv.txt": {
        "abstract": "The formula for the initial mass spectrum of primordial black holes (PBHs), which can be  used for a general case of the scale dependent spectral index, and for a wide class of models of the gravitational collapse, is derived. The  derivation is based on the Press and Schechter formalism. The comparative analysis of different types of initial mass spectra used in concrete calculations is carried out. It is shown that densities of background radiations ($\\nu$, $\\gamma$) from PBH evaporations  depend rather strongly on a type of the gravitational collapse and on a taking into account the spread of horizon masses at which PBHs can form. Constraints on parameters of the primordial density perturbation amplitudes based on PBH evaporation processes and on atmospheric and solar neutrino data are obtained. ",
        "introduction": "Studies of cosmological and astrophysical effects of primordial black holes (PBHs) are important because they enable one to constrain the spectrum of density fluctuations in the early Universe. If the PBHs form directly from primordial density fluctuations then they provide a sensitive probe of the primordial power spectrum on small scales, $\\agt 10^{-9} pc$. In particular, limits on PBHs production can be used to constrain models of inflation, in which the perturbation amplitudes are relatively large at small and medium scales.  In the simplest case, if we assume that the cosmological PBH formation is dominated, approximately, by primordial perturbations of one particular scale (i.e., there is some characteristic epoch of the PBH formation), we can obtain limits on the initial mass fraction of PBHs, \\begin{equation} \\label{int_1} \\rho_i=\\rho_{PBH, i}/\\rho_{tot, i}, \\end{equation} where $\\rho_{PBH, i}$ and $\\rho_{tot, i}$ are the PBH and total energy densities, respectively,  at the time $t_i$ of the formation. This fraction can be expressed by the integral \\begin{equation} \\label{int_2} \\beta_i = \\int\\limits_{\\delta_c}^{1} P(\\delta)d\\delta , \\end{equation} where $P(\\delta)$ is a probability distribution for density fluctuations entering horizon at $t_i$, $\\delta$ is a density contrast, and $\\delta_c$ is a minimum value of $\\delta$ required for the collapse. If the probability distribution  is assumed to be  Gaussian, one has \\begin{equation} \\label{int_3} P(\\delta)=\\frac{1}{\\sqrt{2\\pi}\\sigma} e^{-\\frac{\\delta^2}{2\\sigma^2}}\\:, \\end{equation} where $\\sigma$ is the {\\it rms} fluctuation amplitude on a given scale. It is just this value  that is determined by the primordial power spectrum. The connection between $\\beta$ and $\\sigma$ is very simple in a case of the Gaussian distribution: \\begin{equation} \\label{int_4} \\beta_i\\approx \\sigma e^{-\\frac{\\delta_c^2}{2\\sigma^2}}. \\end{equation}  The limits on $\\beta_i$ arise from the entropy production constraints \\cite{1}, from a distortion of the microwave background \\cite{2}, from the cosmological nucleosynthesis constraints \\cite{3}. In these cases PBHs which give constraints have evaporated completely to the present time. On the contrary, the gravitational constraint ($\\Omega_{PBH, 0}=\\rho_{PBH, 0}/\\rho_c < 1$), \\begin{equation} \\label{int_5} \\beta_i < 10^{-19}\\left(\\frac{M}{10^{15}g}\\right)^{1/2}\\; \\end{equation} \\cite{4}, is valid for PBH masses $M\\agt 10^{15} g$ which survive today.  One should stress that all these limits are based just on the approximation that all PBHs form at the same scale and, correspondingly, the initial mass spectrum of PBHs is $\\delta$-function-like or, at least,  is \"non-extended\" one. Only in this case one can  approximately express the observational constraints through the initial mass fraction $\\beta_i$. Evidently,  a more accurate treatment should operate with the initial PBH mass spectrum directly.  The most strong limit on a PBH formation in  the early Universe is due to the possible contribution of evaporating PBHs to the extragalactic $\\gamma$-ray and neutrino backgrounds at energies $\\sim 100\\text{ MeV}$. The limit  of such kind  was obtained in the work of Page and Hawking \\cite{5}, authors of which assumed that the differential initial mass spectrum of PBHs has power law form predicted in Carr's work \\cite{6},  \\begin{equation} \\label{int_6} n_{BH}(M_{BH})=(\\alpha -2)\\left(\\frac{M_{BH}}{M_{*}}\\right)^{-\\alpha}M_{*}^{-2}\\rho_c\\Omega_{PBH, 0} \\;, \\end{equation} $$ M_{*}\\alt M_{BH}. $$ Here, $\\alpha=2.5$ for PBHs formed in radiation-dominated era, $M_{*}$ is the mass of a black hole whose life-time is equal to the present age of the Universe.  This PBH mass spectrum is, clearly, the example of  an \"extended\" spectrum (it was derived in \\cite{6} by considering the PBH formation as a process stretched in time). Naturally, the famous Page-Hawking constraint was formulated in terms of a PBH number density rather than  in terms of $\\beta_i$. According to \\cite{5}, the upper limit on the present PBH number density is $\\sim 10^{4}pc^{-3}$ (or, that  is the same, $\\Omega_{PBH, 0}\\alt 10^{-8}$). This constraint was improved in later works \\cite{7}, where the same initial PBH mass spectrum was used as input.  The derivation of Eq.(\\ref{int_6}) was based on the assumption of exact scale invariance (scale-independence of the perturbation amplitudes). The Carr's work \\cite{6} appeared before an advent of the inflation hypothesis, and at that time it seemed improbable that the possible case of a growth of the perturbation amplitudes with a decrease of the scale can be of any importance (from the point of view of observational evidences of the PBHs existence). Now we know that, due to inflation, minimum values of PBH mass in the initial mass spectrum can be rather large ($10^{13}-10^{14}$~g or even more). It means that the PBHs can have  rather extended mass spectrum even if the primordial fluctuations are not strictly scale invariant.  In general, initial PBH mass spectrum depends on cosmological and astrophysical aspects of the model used for its derivation (and, through the model, on such parameters as a time of the end of inflation, a reheating temperature, a spectral index of the density perturbations (or parameters of the inflationary potential) and, last but not least,on  parameters, characterizing the process of the gravitational collapse leading to the PBH's birth). Evidently, observational constraints on the PBH production (especially those derived from measurements of extragalactic diffuse backgrounds of $\\gamma$-rays and neutrino) can be used as constraints on at least one of these parameters. Clearly, these constraints  depend on the initial PBH mass spectrum, i.e., on the parameters of this spectrum which are considered as free in a course of the constraint's derivation.  Usually, the parameters characterizing primordial density fluctuation amplitudes (e.g., the spectral index, if it  is scale independent) are objects of the constraining. Parameters of the gravitational  collapse and a method used for a summation over epochs of the PBH  formation (if such a summation is performed) determine, by definition, the {\\it  type} of the  initial PBH  mass spectrum. All  other parameters (characterizing, in particular, the end of inflation and the beginning of radiation dominated era) are \"external\" and  considered as free  ones.  In the present work we obtain some cosmological constraints following from the possible contribution of PBH evaporations in extragalactic diffuse neutrino background. Assuming that the power spectrum of primordial fluctuations has a power-law form, $P(k)\\sim k^n$, we present the observational limits as constraints on values of the spectral index $n$ (using the normalization on COBE data). In a more general case of the non-power $P(k)$-dependence one can directly constrain parameters of a concrete inflation model. We consider, as an example, a case of the running mass inflation model \\cite{8} and obtain constraints on parameters of the corresponding inflationary potential.  The main physical assumption used in the work is that a PBH formation process occurs during radiation dominated era only. The main \"external\" parameter is $t_i$, a time of the beginning of radiation era. At this moment of time we have, by assumption, the primordial spectrum of density perturbations, and no primordial black holes. In a short period between the end of inflation and $t_i$ there can be additional amplification of density perturbation amplitudes (e.g., in a preheating phase) but, by assumption, in this period there is no PBH formation. The time $t_i$ is connected by a usual way with a corresponding initial temperature of radiation era,  which we call a reheating temperature $T_{RH}$.  The paper is organized as follows. In Sec.\\ref{sec:Init_mass_sp} we derive, using the Press and Schechter formalism, the general formula for the initial PBH spectrum which is valid for any law of $P(k)$-dependence and for a wide class of models of the gravitational collapse. In Sec.\\ref{sec:Diff_typ_sp} we give  the comparative analysis of different types of initial PBH mass spectra used in concrete calculations. In Sec.\\ref{sec:Neut_bg_sp} the neutrino background spectra from PBH evaporations are derived and the corresponding constraints on cosmological parameters are obtained. Discussions and conclusions are presented in Sec.\\ref{sec:Dis_and_Con}. ",
        "conclusions": "\\label{sec:Dis_and_Con}  The main conclusion of  the work is the following: constraints on parameters of cosmological models from evaporations of primordial black holes depend rather  strongly on a form of their initial mass spectrum. Number densities of PBHs predicted by the models  are extremely sensitive to values of primordial density perturbation amplitudes and, therefore, even very steeply falling mass spectra of PBHs can give useful constraints. As it follows from Fig.\\ref{fig:fig6}, constraints on the scale independent spectral index $n$ at high values of a reheating temperature are drastically different  in NJ and BK cases. We thoroughly traced the origin of this result: the absence of a  long {\\it high  mass tail} in the initial NJ spectrum (Fig.\\ref{fig:fig1}) leads to a relative shrinkage of the redshift distributions  of evaporated neutrinos (Fig.\\ref{fig:fig2}) and,  as a consequence, to a steepening of the neutrino background spectra and to a decrease of the background intensities (Fig.\\ref{fig:fig4}), resulting, eventually, in a weakening of the spectral index constraint. So, the constraint values at $T_{RH}> 10^{10}\\text{GeV}$ on Fig.\\ref{fig:fig6} are entirely due to  an effect of the summation over all epochs  of the PBH  formation in the approach  of ref.\\cite{20,21}. Analogously,  we see,  on an example of the parameter $c$ of running mass inflation models, that the constraints  at low values of $T_{RH}$ also depend  on a type of the gravitational collapse: in the standard collapse case the constraint at $T_{RH}< 10^{8}$ is practically absent (dotted  curve on Fig.\\ref{fig:fig7}) due  to a corresponding  absence of a {\\it low mass tail} in the initial PBH mass spectrum of this case (compare KL and BK spectra on Fig.\\ref{fig:fig1}; the behavior of the corresponding PBH mass spectra in a case of a scale dependent $n$ is qualitatively similar).  The initial number density of PBHs is larger in a case of the standard picture of the gravitational collapse (as compared with the critical collapse case). It is due to relatively small value of the critical overdensity ($\\delta_c^{st}=\\gamma=1/3$). In the region near the maximum of PBH mass spectra the ratio of intensities is given by the approximate relation \\begin{equation} \\label{12} \\frac{n_{BH}^{(KL)}}{n_{BH}^{(BK)}}\\sim e^{\\frac{\\delta_c^2-\\gamma^2}{2\\sigma_H^2}}\\;. \\end{equation} Correspondingly, intensities of the neutrino and gamma backgrounds produced by PBH's evaporations are smaller and spectral index constraints are systematically weaker in the critical collapse case (see, for  the comparison with  the standard collapse case, Fig.10 of ref.\\cite{11}).  It is worth noting that spectral index constraints followed from neutrino evaporations are stronger at high $T_{RH}$ values (as compared with the constraints from $\\gamma$-quanta), especially in the NJ case (Fig.\\ref{fig:fig6}). The reason is simple: at high $T_{RH}$ the relative contribution of large redshifts in $z$-distributions is very large (Fig.\\ref{fig:fig2}) and, correspondingly, the absorption of $\\nu$,~$\\gamma$ during propagation in space is important. Naturally, absorption effects are more sufficient for $\\gamma$-quanta than for neutrinos (for $\\gamma$-quanta $z_{max}\\sim 700$, while for neutrinos $z_{max}\\sim 10^7$ \\cite{11}, as we can see, in particular, from curves  of Fig.\\ref{fig:fig2}).  One should  note that, in general, the constraints on a scale independent spectral  index followed from  PBH evaporations (Fig.\\ref{fig:fig6})  are  rather  weak. Probably, so large deviations of the spectral index from 1 ($n\\sim 1.28$) are excluded by the latest data (see,  e.g., \\cite{30}). Therefore, constraints from PBH evaporations  are  more useful for those  cosmological models in which the spectral  index is  scale dependent. In such models large density perturbation amplitude at small scales (and, correspondingly, large effects from evaporations) can  coexist with the small value of    $\\delta_H (k)$ at COBE scale (Fig.\\ref{fig:fig8}). We  showed in this paper, using, as an example, the running mass inflation  model, that in this case  the PBH evaporation  process can  be used for  a  constraining of model parameters [in particular, parameters of an inflationary potential, (Fig\\ref{fig:fig7})]. We stress, once more, that all constraints of such type depend on assumptions used in a derivation of the initial PBH mass spectrum."
    },
    "0206/astro-ph0206031_arXiv.txt": {
        "abstract": "{This paper presents preliminary results of a spectroscopic survey being conducted at the VLT of fields with optically-selected cluster candidates identified in the EIS $I$-band survey. Here we report our findings for three candidates selected for having estimated redshifts in the range $z=0.8-1.1$. New multi-band optical/infrared data were used to assign photometric redshifts to galaxies in the cluster fields and to select possible cluster members in preparation of the spectroscopic observations. Based on the available spectroscopic data, which includes 147 new redshifts for galaxies with $I_{AB}\\lsim22-23$, we confirm the detection of four density enhancements at a confidence level $>99\\%$.  The detected concentrations include systems with redshifts $z=0.81$, $z=0.95$, $z=1.14$ and the discovery of the first optically-selected cluster at $z=1.3$.  The latter system, with three concordant redshifts, coincides remarkably well with the location of a firm X-ray detection ($>5\\sigma$) in a $\\sim80$~ksec XMM-Newton image taken as part of this program which will be presented in a future paper (Neumann \\etal 2002). The $z>1$ systems presented here are possibly the most distant identified so far by their optical properties alone. ",
        "introduction": "\\label{sec:intro}  Clusters of galaxies are both ideal sites for studying galaxy evolution and important cosmological probes, especially at redshifts $z\\gsim0.5$, where differences between competing evolutionary and cosmological models become important. This has motivated several searches for distant clusters using a variety of techniques in different wavelengths. As a result, over the past few years a remarkable progress has been made in detecting an ever increasing number of systems with $z \\gsim0.5$ (see Gioia 2000 for a recent review). More recently, a handful of clusters at $z\\gsim1$ have also been identified. While the sheer existence of these high redshift clusters is of great importance, the current number of confirmed systems is still very small, mostly identified from serendipitous X-ray searches (\\eg Rosati \\etal 1999) or from infra-red imaging (Stanford \\etal 1997).  Therefore, the construction of a large sample of confirmed clusters at $z\\gsim0.8$ representative of the entire population of these high-z systems remains an important goal of observational cosmology. However, as these systems are expected to be rare, finding them requires large areas of the sky to be covered, limiting the techniques that can be used in identifying candidates. In particular, surveys at X-ray and mm wavelengths (Carlstrom \\etal 2000) are unlikely to provide in the near future the necessary sky coverage for constructing the large samples of very distant clusters of galaxies required for statistical analyses.  An alternative way is to consider multi-band optical/infrared imaging data. Thanks to the advent of panoramic CCD imagers, wide-angle imaging surveys in the optical and near infrared wavelengths have become viable and can be used for identifying cluster candidates up to $z\\sim1$. Examples of wide-angle surveys that have been used to identify intermediate to high redshift clusters include those of Gunn \\etal (1986), Postman \\etal (1996), the ESO Imaging Survey (EIS) Cluster Survey (Olsen \\etal 1999a,b; Scodeggio \\etal 1999), the Red-Sequence Survey (Gladders \\& Yee 2000) and the Las Campanas distant cluster survey (Gonzalez \\etal 2001).  These surveys, especially those carried out in a single passband, can only provide plausible candidates and further benefits from additional multi-wavelength observations to mitigate many of the problems of foreground-background contamination, to assign photometric redshifts for galaxies of different morphological types and to select possible cluster members to improve the yield of spectroscopic follow-ups.  In this paper we describe our first attempts to explore the nature of the high redshift cluster candidates identified in the EIS $I$-band survey, combining new imaging and spectroscopic observations. Altogether there are about 82 candidates with matched-filter redshifts $\\gsim 0.8$ for which about half have already been complemented by imaging observations in $BVRJK$. Among the various clusters for which we have spectroscopic data, we present here three clusters at higher redshift.  In section~\\ref{sec:data} we describe the selection of the candidate clusters and of the galaxy sample used in the observations. In section~\\ref{sec:obs_red}, we briefly describe the reduction procedure which will be expanded in a separate paper where the accumulated data are presented (J{\\o}rgensen \\etal 2002). In section~\\ref{results}, the observed redshift distribution and the technique used to identify groups in redshift space are presented. Finally, in section~\\ref{summary} our main results are summarised. ",
        "conclusions": "\\label{summary}    This paper presents new spectroscopic data of EIS cluster candidate fields identified from moderately deep $I$-band images using the matched-filter algorithm. The three fields considered were selected because the cluster candidates had estimated redshifts beyond $z=0.8$. Analysis of the spectroscopic data strongly suggests the existence of real density enhancements at high redshifts ($0.8<z<1.3$) in all of them. The measured redshifts are, in general, consistent with those estimated from the photometric data. In at least one of the cluster fields, the location of the high-redshift system coincides remarkably well with a robust X-ray detection, lending further support to the reality of the system.  Therefore, despite the fact that it is difficult to decide if those clumps are filaments or bound systems, for two of them the evidence seem to favour the second possibility: for the $z=0.81$ clump a red sequence is observed (da Costa \\etal 1999) whereas the $z=1.3$ one has an X-ray detection associated with it.  The present paper strongly suggests that the EIS cluster candidate catalog provides a valuable pool from which to construct a statistical sample of optically-selected clusters at high redshift. The present data alone contribute with four systems at $z>0.8$ in the southern hemisphere, two of which at $z\\gsim1$, ideal for VLT studies. The success in identifying significant concentrations from a relative small sample underscores the importance of collecting multi-band optical/infrared data and estimating photometric redshifts to select potential cluster members. However, in establishing the true nature of these systems will require a better sampling of these systems which will become possible with the availability of an integral field unit as foreseen by the VIMOS spectrograph."
    },
    "0206/astro-ph0206207_arXiv.txt": {
        "abstract": "Recent spectroscopic and high resolution $HST$-imaging observations have revealed significant numbers of ``passive'' spiral galaxies in distant clusters, with all the morphological hallmarks of a spiral galaxy (in particular, spiral arm structure), but with weak or absent star formation. Exactly how such spiral galaxies formed and whether they are the progenitors of present-day S0 galaxies is unclear. Based on analytic arguments and numerical simulations of the hydrodynamical evolution of a spiral galaxy's halo gas (which is a likely candidate for the source of gas replenishment for star formation in spirals), we show that the origin of passive spirals may well be associated with halo gas stripping. Such stripping results mainly from the hydrodynamical interaction between the halo gas and the hot intracluster gas. Our numerical simulations demonstrate that even if a spiral orbits a cluster with a pericenter distance $\\sim$  3 times larger than the cluster core radius, $\\sim$  80 \\% of the halo gas is stripped within a few Gyr and, accordingly, cannot be accreted by the spiral. Furthermore, our study demonstrates that this dramatic decline in the gaseous infall rate leads to a steady increase in the $Q$ parameter for the disk, with the spiral arm structure, although persisting, becoming less pronounced as the star formation rate gradually decreases. These results suggest that passive spirals formed in this way, gradually evolve into red cluster S0s. ",
        "introduction": "One of remarkable conclusions of observational studies of galaxy evolution in distant (z $>$ 0.2) clusters is that galaxies are undergoing both spectrophotometric and morphological evolution, but the timescales appear to be different for each (Couch et al. 1998; Dressler et al. 1999; Poggianti et al. 1999). In particular, these studies found a significant number of galaxies which had the morphological appearance of spiral galaxies, but which had no evidence of ongoing or even recent star formation in their spectra. The existence of these so-called ``passive'' spirals (or ``k-type'' spirals; Dressler et al. 1999), coupled with the strong evidence for a morphological transformation of spirals into S0s in rich clusters since $z\\sim 0.5$ (Dressler et al. 1997), suggests that there are two different time scales in cluster galaxy evolution. One is the time-scale for the suppression of star formation that eventually leads to spirals having passive, k-type spectra. The other, which clearly is much longer (Poggianti et al. 1999), is the time-scale for the morphological transformation of a spiral into an S0. However, it remains unclear as to (i)\\,why the morphological transformation is preceded by spectral evolution, (ii)\\,what cluster-related physical processes are responsible for passive spiral formation, (iii)\\,how or whether passive spirals might finally be transformed into passive S0s, and (iv)\\, whether  there is an evolutionary link between the so-called ``E+A'' galaxies discovered by Dressler \\& Gunn (1983, 1992) and passive spirals.  In one of the first theoretical attempts to explain the `Butcher-Oemler' effect -- which first drew attention to these evolutionary effects going on in rich clusters (Butcher \\& Oemler 1978) -- Larson, Tinsley, \\& Caldwell (1980) suggested that infall from gaseous halos might be important for sustaining star formation in spirals, and that these halos might be stripped in the cluster environment, leading to the formation of S0s. Not only might this explain the demise of the blue, Butcher-Oemler populations in distant clusters, but also the presence of the smooth-armed `anemic' spirals identified by van den Bergh (1976) in present-day clusters. More generally, replenishment of interstellar gas due to sporadic and continuous gas infall and acquisition from external environments has also been considered to provide reasonable and plausible explanations for the problems of gas consumption time-scales for the Galaxy and typical late-type galaxies (Kennicutt 1983), the G-dwarf problem in the Galaxy (e.g., van den Bergh 1962), and the formation of counter-rotating components in disk galaxies (e.g., Bettoni, Galletta, \\& Oosterloo 1991; Bertola, Buson, \\& Zeilinger 1992). Furthermore, gas accretion from extended diffuse halo gas due to radiative cooling is critically important for disk galaxy formation in a hierarchical clustering scenario (White \\& Frenk 1991), although the observed diffuse X-ray emission from late-type spirals appears to be inconsistent with the predictions from this scenario (Benson et al. 2000).  While Larson et al's halo gas stripping scenario provides a very interesting possibility for transforming spirals into S0s, little detailed work has been done to understand how passive spirals fit within this framework. The purpose of this paper is to investigate whether the hydrodynamical interaction between a gaseous halo that might surround a spiral galaxy and the hot intracluster gas is likely to lead to the formation of a passive spiral. We do this both analytically and numerically, examining the effects the intracluster gas has on halo gas reservoirs of spirals orbiting a cluster. In particular, we determine how the total amount of halo gas stripped as a result of this interaction, depends on both the orbit of the galaxy within the cluster and the mass of the cluster. We also present numerical results on the morphological evolution of spiral galaxies whose gas infall rates are declining as a result of halo gas stripping. Based on these numerical results, we discuss the possible transformation of blue spirals into red S0s via a passive spiral phase. ",
        "conclusions": "We have investigated, numerically, the dynamical interaction between the halo gas of spiral galaxies and the hot ICM in order to estimate the amount of halo gas that can be stripped from spirals which reside in rich clusters. We confirmed that this hydrodynamical interaction is an efficient means of stripping the halo gas in cluster spiral galaxies. We have also investigated the dynamical evolution of spiral galaxies that have a very small gas accretion rate due to the removal of their halo gas. We demonstrated that the spiral arm structure in these disk galaxies becomes rapidly less pronounced and then eventually disappears. Coupled with this is a dramatic drop in the star formation rate due to the rapid increase in the $Q$ parameter -- a key quantity in determining the occurrance of star formation in disk galaxies. We therefore conclude that halo gas stripping caused by dynamical interaction between halo gas and the hot ICM is a plausible mechanism for not just S0 production in distant clusters but also the passive spirals observed in distant clusters    This present study also demonstrated that the dramatic tidal effects associated with galaxy-galaxy interactions and merging and the mass distribution of the cluster itself are not necessarily the sole cause of environmental differences in galaxy evolution between rich clusters and the field. Passive spirals recently discovered in distant clusters (Couch et al. 1998) and anemic spirals and late-type disk galaxies with low levels of star formation observed in nearby groups and clusters (van den Bergh 1976, Tran et al. 2001) are probably objects which provide valuable information on less dramatic but long-term environmental effects on galaxy evolution. We lastly suggest that formation of some cluster S0s  can be due to this long-term environmental effects."
    },
    "0206/astro-ph0206494_arXiv.txt": {
        "abstract": "Microlensing is sensitive to binary, brown dwarf (BD), and planetary companions to normal stars in the Galactic bulge with separations between about 1-10 AU.  The accurate, densely-sampled photometry of microlensing events needed to detect planetary companions has been achieved by several follow-up collaborations.  Detailed analysis of microlensing events toward the bulge demonstrates that less than 45\\% of M-dwarfs in the bulge have $\\mjup$ companions between 1 and 5 AU.  Detection of binary and BD companions using microlensing is considerably easier; however, the interpretation is hampered by their non-perturbative influence on the parent lightcurve.  I demonstrate that $\\sim 25\\%$ of BD companions with separations $1-10{\\rm AU}$ should be detectable with survey-quality data ($\\sim 1~{\\rm day}$ sampling and $\\sim 5\\%$ photometry).  Survey data is more amenable to generic, brute-force analysis methods and less prone to selection biases.  An analysis of the $\\sim 1500$ microlensing events detected by OGLE-III in the next three years should test whether the BD desert exists at separations $1-10{\\rm AU}$ from M-dwarfs in the Galactic bulge. ",
        "introduction": "It is by now well-established by radial velocity (RV) surveys that the mass function of low-mass companions to normal (GKM) stars in the solar neighborhood exhibits a ``brown-dwarf (BD) desert,'' a paucity of $13-80\\mjup$ companions.  Specifically, $<1\\%$ of solar-type stars have BD companions with semi-major axes $a\\la 5\\au$ (Marcy \\& Butler 2000).  Although RV surveys will eventually be able to detect more distant companions, the duration of their observations are currently only sufficient to detect long-term trends from companions with $a\\ga 5\\au$ (See Figure 1). Young BD companions with separations greater than a few tens of AUs and less than a few hundreds of AUs can be detected via direct imaging surveys (Oppenheimer et~al.\\ 2001, Lowrance 2001, McCarthy 2001); even more distant companions are detectable serendipitously in wide-field surveys (Gizis et~al.\\ 2001).  These studies seem to suggest that the frequency of wide BD companions to GKM stars is $\\la 10\\%$ (see McCarthy 2001).  Neither RV nor direct imaging can currently assess the frequency of companions with $a=5-10\\au$.  The exquisite precision ($\\sim 3~\\kms$) of RV surveys also makes them sensitive to planetary companions with mass $0.1-13\\mjup$; currently nearly 100 such companions are now known.  The $M\\sin i$ and $a$ of these companions are shown in Figure 1.  There also appears to be a statistically significant dearth of high-mass ($M>5\\mjup$), close-in ($a\\la 0.3\\au$) planetary companions (Zucker \\& Mazeh 2002; see also Figure 1).  Whether this is related to the BD desert is not clear.  Microlensing is sensitive to binary, BD, and planetary companions separated by $\\sim 1-10\\au$ from the objects in the Galactic bulge which serve as the lenses of microlensing events detected toward the Galactic bulge (Mao \\& Paczy\\'nski 1991).  The majority of the primaries are likely to be M-dwarfs in the bulge, with some contamination from stellar remnants and disk stars.  Microlensing is sensitive to the mass ratio $q$ and projected separation $d$ of binary systems, where $d$ is in units of the combined Einstein ring radius $r_{\\rm E}$ of the system.  I will assume that the primaries are M-dwarfs with $M=0.3\\msun$ in the bulge, so that $r_{\\rm E}=2\\au$, and use these to transform to the mass $M_p$ and semi-major axis $a$ of the secondaries. ",
        "conclusions": ""
    },
    "0206/astro-ph0206177_arXiv.txt": {
        "abstract": "We report the results of [O~III] $\\lambda 5007$ surveys for planetary nebulae (PNe) in six galaxies: NGC~2403, NGC~3115, NGC~3351, NGC~3627, NGC~4258, and NGC~5866.  Using on-band/off-band [O~III] $\\lambda 5007$ images, as well as images taken in H$\\alpha$, we identify samples of PNe in these galaxies and derive their distances using the planetary nebula luminosity function (PNLF).  We then combine these measurements with previous data to compare the PNLF, Cepheid, and surface brightness fluctuation (SBF) distance scales.  We use a sample of 13 galaxies to show that the absolute magnitude of the PNLF cutoff is fainter in small, low-metallicity systems, but the trend is well modeled by the theoretical relation of \\citet{djv92}.  When this metallicity dependence is removed, the scatter between the Cepheid and PNLF distances becomes consistent with the internal errors of the methods and independent of any obvious galaxy parameter.  We then use these data to recalibrate the zero point of the PNLF distance scale.  We use a sample of 28 galaxies to show that the scatter between the PNLF and SBF distance measurements agrees with that predicted from the techniques' internal errors, and that there is no systematic trend between the distance residuals and stellar population.  However, we also find that the PNLF and SBF methods have a significant scale offset:  Cepheid-calibrated PNLF distances are, on average, $\\sim 0.3$~mag smaller than Cepheid-calibrated SBF distances.  We discuss the possible causes of this offset, and suggest that internal extinction in the bulges of the SBF calibration galaxies is the principle cause of the discrepancy.   If this hypothesis is correct, then the SBF-based Hubble Constant must be increased by $\\sim 7\\%$.  We also use our distance to NGC~4258 to argue that the short distance scale to the LMC is correct, and and that the global Hubble Constant inferred from the {\\sl HST\\/} Key Project should be increased by $8 \\pm 3\\%$ to $H_0 =  78 \\pm 7$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "introduction": "The past two decades has seen remarkable progress in the measurement of the distance scale of the universe.  In the early and mid-1980's, values of the Hubble Constant ranged over a factor of two, from $H_0 \\approx 50$~km~s$^{-1}$~Mpc$^{-1}$ \\citep[\\eg][]{kkct88, st82} to $H_0 \\approx 100$~km~s$^{-1}$~Mpc$^{-1}$ \\citep[\\eg][]{deV85, huc87}, with the results depending strongly on the the author and the technique. However, in the early 1990's, measurements of $H_0$ began to converge \\citep{mudville}, and it became increasingly difficult to argue for values much different than $H_0 \\sim 70$~km~s$^{-1}$~Mpc$^{-1}$.  Today, due in large part to the {\\sl Hubble Space Telescope\\/} Distance Scale Key Project, a value of $H_0$ between 65 and 75 km~s$^{-1}$~Mpc$^{-1}$ is generally accepted \\citep{keyfinal}.  Nevertheless, there are still two lingering problems with the current distance scale.  The first concerns the zero point of the Cepheid period-luminosity relation.  There are two galaxies whose distances are known from direct geometric techniques: the Large Magellanic Cloud \\citep[via the light echo of SN~1987A;][]{panagia91, gould}, and NGC~4258 \\citep[through the observed motions of its nuclear maser;][]{herrnstein}.  The former sets the zero point for the Cepheid scale; the latter provides an independent test of the technique.  Unfortunately, the Cepheid distance to NGC~4258 given by \\citet{keyfinal} is $1.2 \\sigma$ larger than the galaxy's geometric distance \\citep{herrnstein}.  This marginally significant discrepancy may indicate a problem for the zero point of the system.  The second limitation of the Cepheid distance scale is its limited applicability to Population~II distance techniques.  For example, the calibration of the elliptical galaxy fundamental plane \\citep{kelson, keyfinal} rests largely on the assumption that the early-type galaxies of Leo~I, Fornax, and Virgo are at the same distance as the clusters' spirals.  In the case of Virgo at least, this is likely not the case \\citep{cm87, wb00}.  Similarly, the zero point of the surface brightness fluctuation (SBF) technique \\citep{ferrarese, tonry, keyfinal} is set by just six Cepheid calibrators. Given the susceptibility of the SBF method to the effects of interstellar extinction, this situation is not ideal.  Clearly, additional calibrators are needed to secure the Pop~II side of the distance ladder.  The planetary nebula luminosity function (PNLF) has the potential to provide these calibrations.  As the only general purpose standard candle that is applicable to both spiral and elliptical galaxies, the PNLF provides a critical link between the Pop~I and Pop~II distance scales.   Moreover, since the precision of the PNLF method is comparable to that of Cepheids \\citep{mudville}, the technique can also be used to check for anomalous measurements in the distance ladder.  In fact at present, the PNLF is the only method capable of confirming the results obtained from Cepheid variables. Finally, the PNLF can provide distances to some intermediate objects that are too dusty or irregular for Pop~II techniques, but not suitable for Cepheid observations.  In this paper, we present the PNLFs of six galaxies, NGC~2403, 3115, 3351, 3627, 4258, and 5866, and use these data to search for systematic errors in the extragalactic distance ladder.  In Section 2, we describe our observations, detail our reduction procedures, and present the coordinates and [O~III] $\\lambda 5007$ magnitudes of our planetary nebula candidates.  We also present new [O~III] and \\Halpha observations of PNe in the inner bulge of M31; these data are used in Section~3 to create a quantitative criterion for discriminating PNe from compact H~II regions.  In Section~4, we derive PNLF distances to our six galaxies, and comment on the properties of these systems.  Included in this section is a discussion of the distance to NGC~4258; our value, combined with that obtained from the Cepheids, argues for a Hubble Constant that is $\\sim 7\\%$ larger than that given by the {\\sl HST\\/} Key Project \\citep{keyfinal}.  In Section~5, we combine our distances to NGC~2403, 3351, 3627, and 4258 with data from nine other Cepheid galaxies to re-define the zero point of PNLF distance scale. We show that the absolute magnitude of the PNLF bright-end cutoff does shift to fainter magnitudes at extremely low metallicity;  this is in agreement agreement with the theoretical predictions of \\citet{djv92}.  However, we show that in metal-rich galaxies, the PNLF-Cepheid residuals show no statistically significant trend.  In Section~6, we compare the PNLF distance scale with that of the SBF method, and show that there is a significant scale error between the two techniques.  Specifically, we show that, although the PNLF-SBF residuals do not correlate with any galaxy property, the overall PNLF scale is $\\sim 0.3$~mag shorter than the SBF scale.  Finally, we conclude by considering the possible causes of this discrepancy, and discussing the implications it has for tip of the red giant branch (TRGB) distance measurements and the extragalactic distance scale in general. ",
        "conclusions": "We have presented PNLF distances to NGC~2403, NGC~3115, NGC~3351, NGC~3627, NGC~4258, and NGC~5866, and have used these data to compare the Cepheid, PNLF, and SBF distance scales.  Our observations demonstrate that, in terms of relative distance measurements, the Cepheid, PNLF, and SBF methods are in excellent agreement, and the internal errors estimated for all the methods are correct.  However, we also show that the PNLF and SBF distance scales are incompatible:  the Cepheid-calibrated SBF scale is $\\sim 0.3$~mag longer than the Cepheid-calibrated PNLF scale.  The likely cause of the discrepancy is internal extinction in the bulges of the SBF Cepheid calibrators.  If this is true, then this error results in an underestimate of the SBF Hubble Constant.  Finally, we use our PNLF distance to NGC~4258, in combination with the galaxy's geometric and Cepheid distances, to argue that the short distance to the Large Magellanic Cloud is correct, and that the \\citet{keyfinal} Hubble Constant should be increased by $8 \\pm 3\\%$."
    },
    "0206/astro-ph0206388_arXiv.txt": {
        "abstract": "The halo approach to large scale structure provides a physically motivated model to understand clustering properties of galaxies. An important aspect of the halo model involves a description on how galaxies populate dark matter halos or what is now called the halo occupation distribution. We discuss a way in which clustering information, especially in the non-linear regime, can be used to determine moments of this halo occupation number. We invert the non-linear part of the real space power spectrum from the PSCz survey  to determine the second moment of the halo occupation distribution in a model independent manner.  The precise measurement of higher order correlations can eventually be used  to determine successive higher order moments of this distribution. ",
        "introduction": "The halo approach to large scale structure has now become a useful tool to study and understand  clustering properties of dark matter and a number of tracers including galaxies (see \\cite{CooShe02} 2002 for a recent review). This approach replaces the complex distribution of dark matter with a collection of collapsed dark matter halos. Thus, necessary inputs for a halo based model include properties of this dark matter halo population, such as its mass function and the spatial profile of dark matter within each halo (\\cite{Sel00} 2000; \\cite{MaFry00} 2000; \\cite{Scoetal00} 2000; \\cite{Cooetal00} 2000).  In order to describe clustering aspects beyond dark matter, it is necessary that one understands how the tracer property is related to the dark matter distribution in each halo. In the case of galaxies, an important input is a description on how galaxies populate halos. This is usually achieved by the so-called halo occupation number where one describes the mean number of galaxies in dark matter halos as a function of mass and its higher order moments (\\cite{Sel00} 2000; \\cite{PeaSmi00} 2000; \\cite{BerWei02} 2002). For the two-point correlation function, one requires information up to the second moment of the halo occupation distribution, while higher order correlations successively depend on  increasing moments.  The halo occupation distribution has been widely discussed in the literature in terms of semi-analytical models of galaxy formation (e.g., Benson et al. 2001; Somerville et al. 2001). With the advent of a well defined halo approach to clustering, observational constraints have also begun to appear (e.g., \\cite{Scoetal00} 2000; \\cite{MouSom02} 2002). While expectations for constraints on the halo occupation distribution from current wide-field galaxy surveys, such as the Sloan Digital Sky Survey, are high (\\cite{BerWei02} 2002; \\cite{Scr02} 2002), these are, however, all considered in a model dependent manner.  Though descriptions on the halo occupation number based on a specific model is useful, it is probably more useful to consider model independent constraints. In this {\\it Letter}, we consider such an approach and suggest that clustering information, especially in the non-linear regime, can be used for a reconstruction of various moments of the halo occupation number. We discuss a possible inversion for this purpose and use results on the non-linear power spectrum from the PSCz redshift survey (Saunders et al. 2000) by Hamilton \\& Tegmark (2002; see also, Hamilton et al. 2000) to provide a first estimate of the second moment of the halo occupation number. We discuss both strengths and limitations of our approach and provide a comparison to model based descriptions of the halo occupation number.  We provide a general discussion of our method in the next section. when illustrating results in \\S~3, we take a flat $\\Lambda$CDM cosmology with parameters $\\Omega_c = 0.3$, $\\Omega_b=0.05$, $\\Omega_\\Lambda=0.65$, $h=0.65$, $n=1$, $\\delta_H=4.2\\times 10^{-5}$. ",
        "conclusions": "The halo approach to large scale structure provides a physically motivated technique to study clustering properties of galaxies (Cooray \\& Sheth 2002 and references therein). A necessary, and an important, ingredient  for a halo based clustering calculation involve a description of how galaxies populate dark matter halos or the so-called halo occupation number. We have raised the possibility for a model independent study on moments of the halo occupation number using an inversion of the non-linear clustering power, and p-point, spectrum measurements. We have considered an application of this suggestion utilizing the PSCz power spectrum estimated by  Hamilton \\& Tegmark (2002). Our estimates on the second moment of the halo occupation number are  consistent with power law models  over five decades in mass and with certain model descriptions in the literature. With expected increase in measurements of the clustering in the non-linear regime, and associated measurements of covariance, we expect analysis like the one suggested here will eventually make it possible for a detailed understanding of the nature of galaxy occupation in halos.  \\smallskip {\\it Acknowledgments:} This research was supported at Caltech by a senior research fellowship from the  Sherman Fairchild foundation and a DOE grant (DE-FG03-92-ER40701). We thank Ryan Scranton and Andrew Hamilton for useful suggestions and appreciate the quite atmosphere at the Aspen Center for Physics where this work was initiated."
    },
    "0206/astro-ph0206341_arXiv.txt": {
        "abstract": "General importance and capabilities of observations of eclipsing binaries by the forthcoming ESA mission GAIA are discussed. Availability of spectroscopic observations and a large number of photometric bands on board will make it possible to reliably determine physical parameters for $\\sim 10^5$ binary stars. It is stressed that current methods of object by object analysis will have to be modified and included in an automatic analysis pipeline. ",
        "introduction": "GAIA is the approved Cornerstone 6 mission of the European Space Agency. Its main goal is to observe up to a billion stars in our Galaxy and obtain their astrometric positions on a micro-arc sec level, multi-band photometry in 15 different optical and near-IR bands, as well as spectroscopic observations within the  250~\\AA\\ interval around the Ca~II IR triplet. Perryman et al. (2001), ESA-SCI(2000)4 and Munari (1999, 2001) are useful introductions to the general properties of the mission, its overall astrophysical importance and diagnostic capabilities of the spectral window chosen for the spectrograph. Munari (2002) discusses the potential of GAIA to observe peculiar stars. Here we focus on galactic eclipsing binaries which are rapidly becoming one of the areas where the harvest of GAIA's results will be the richest.  In the next section we discuss general importance and capabilities of observations of eclipsing binaries by GAIA. Next the mode in which spectroscopic observations are obtained is briefly discussed, together with sources of noise, intensity of the background and the possibility for spectral tracing overlaps. All this leads to the estimates on expected radial velocity accuracy as a function of magnitude and spectral type of the target. Finally we present results from observations of real binary stars in the GAIA-like mode and so demonstrate GAIA's accuracy in determination of basic stellar parameters. We conclude with some remarks on automatic reduction of the gigantic database of GAIA's observations of galactic eclipsing binaries. ",
        "conclusions": "The GAIA mission has the potential to discover and characterize thousands of new peculiar and binary stars that could be followed up by additional observations by automatic ground based telescopes. It is crucial that suitable codes are developed for a correct recognition of eclipsing and non-eclipsing binaries on a central reduction pipeline. Once recognized the light curves and spectra will need to be analyzed with an automatic and reliable automatic procedures.  The unique potential of GAIA among the forthcoming space missions is the availability of spectroscopy on board and an unprecedented number of targets that are to be observed. The size of the database at the end of the mission will exceed 10 TB. However, as correctly pointed out by Peter Eggleton, more data can mean less understanding. To avoid this it is very important to think now of the new era when even peculiar binaries will be observed in large numbers. \\medskip  {\\bf Acknowledgment.} This work has been supported by a grant from the Slovenian Ministry of Education, Science and Sport."
    },
    "0206/astro-ph0206406_arXiv.txt": {
        "abstract": "We present K$'$-band observations of five Mira stars with the IOTA interferometer. The interferograms were obtained with the FLUOR fiber optics beam combiner, which provides high-accuracy visibility measurements in spite of time-variable atmospheric conditions. For the M-type Miras \\astrobj{X Oph}, \\astrobj{R Aql}, \\astrobj{RU Her}, \\astrobj{R Ser}, and the C-type Mira \\astrobj{V CrB} we derived the uniform-disk diameters 11.7~mas, 10.9~mas, 8.4~mas, 8.1~mas, and 7.9~mas ($\\pm$0.3 mas), respectively. Simultaneous photo\\-metric observations yielded the bolometric fluxes. The derived angular Rosseland radii and the bolometric fluxes allowed the determination of effective temperatures. For instance, the effective temperature of \\astrobj{R Aql} was determined to be 2970\\,$\\pm$110~K. A linear Rosseland radius for \\astrobj{R Aql} of 250$^{+100}_{-60}$\\,R$_{\\odot}$ was derived from the angular Rosseland radius of 5.5\\,mas\\,$\\pm$0.2\\,mas and the HIPPARCOS parallax of 4.73 mas\\,$\\pm$1.19\\,mas. The observations were compared with theoretical Mira star models of \\citet{bsw96} and \\citet{hsw98}. The effective temperatures of the M-type Miras and the linear radius of \\astrobj{R Aql} indicate fundamental mode pulsation. ",
        "introduction": "The resolution of large optical telescopes and interferometers is high enough to resolve the stellar disk of nearby M giants, to reveal photospheric asymmetries and surface structures, and to study the dependence of the diameter on wavelength and variability phase \\citep[see e.g. pioneering work by][]{bon73,karov91}. Theoretical studies \\citep[e.g.][]{wata79,schol85,bess89,bsw96,hsw98} show that accurate monochromatic diameter measurements can significantly improve our understanding of M giant atmospheres. With the IOTA (= Infrared-Optical Telescope Array) interferometer, a resolution of $\\sim$\\,11.9~mas can be achieved with its largest baseline of 38~m in the K$'$-band. The IOTA interferometer is located at the Smithsonian Institution's Whipple Observatory on Mount Hopkins in Arizona. A detailed description of IOTA can be found in \\citet{carle94} and \\citet{trau98}. IOTA can be operated in the K-band with the FLUOR (= Fiber Linked Unit for Optical Recombination, \\citeauthor{fore97} \\citeyear{fore97}) fiber optics beam combiner. This beam combiner provides high-accuracy visibility measurements in spite of time-variable atmospheric conditions. The single-mode fibers in the beam combiner spatially filter the wavefronts corrugated by atmospheric turbulence \\citep{fore97,perr98}. ",
        "conclusions": "We derived angular uniform-disk diameters $\\Theta_{\\rm UD}$ of five Mira stars (Table~\\ref{tab:obs}) from K$'$-band visibility measurements with the IOTA interferometer and the FLUOR beam combiner at 38\\,m baseline. Using simultaneously observed bolometric fluxes and the measured uniform-disk diameters we obtained T$_{\\rm eff,UD}$ values given in Table~\\ref{tab:Fboltab}.  Previous interferometric K-band observations of some of our target stars (\\astrobj{R Aql}, \\astrobj{X Oph}, \\astrobj{R Ser}) were carried out by \\citet{bell96} at similar phases. Their derived uniform-disk diameters (\\astrobj{R Aql}; $\\Phi_{\\rm vis}$\\,=\\,0.90: 10.76$\\pm$0.61\\,mas, \\astrobj{X Oph}; $\\Phi_{\\rm vis}$\\,=\\,0.75: 12.30$\\pm$0.66\\,mas, \\astrobj{R Ser}; $\\Phi_{\\rm vis}$\\,=\\,0.32: 8.56$\\pm$0.58\\,mas) are in good agreement with our observations. Their effective temperatures, derived from measured angular Rosseland radii (\\astrobj{R Aql}: 3189$\\pm$147\\,K, \\astrobj{X Oph}: 3041$\\pm$160\\,K, \\astrobj{R Ser}: 2804$\\pm$144\\,K), are also in agreement with our results.  The comparison of the observations with Mira star models, with respect to the effective temperature, suggests that the four observed M Miras can approximately be represented by the fundamental mode models D, P or M, whereas the overtone models are much too cool. For the two M stars \\astrobj{R Aql} and \\astrobj{RU Her}, phase- and model-dependent effective temperatures could be derived, which are listed in Table~\\ref{tab:Fboltab}. These effective temperatures are within the error bars of the T$_{\\rm eff}$ values obtained with the measured uniform disk diameter. Any more accurate model interpretation of our four M-type stars would require an extension of the parameter range (Table~\\ref{tab:prop}) and refining the phase spacing (Table~\\ref{tab:link}) of available Mira model grids. A quantitative study of \\astrobj{V CrB} can only be given on the basis of C-type Mira models which are not yet available.  For \\astrobj{R Aql}, a useful HIPPARCOS parallax (4.73$\\pm$1.19\\,mas) is available and it is therefore possible to compare measured linear Rosseland and stellar K$'$-band radii with the theoretical radii of the BSW96 and HSW98 models. The measured radii were derived by fitting theoretical (BSW96, HSW98) center-to-limb intensity variations to the visibility data. In Table~\\ref{tab:Linear}, the measured linear Rosseland radii derived from the well fitting near-maximum fundamental mode models D, P and M are listed. From the measured linear Rosseland radii of \\astrobj{R Aql} the pulsation mode could not be determined because of the large parallax error.  The comparison suggests that \\astrobj{R Aql} can well be represented by the fundamental mode D or P model. Note, however, that observations in more filters than just one continuum filter and more baselines may be necessary for safely distinguishing a well-fitting model from an accidental match (cf. Hofmann et al. 2000a). Furthermore, in order to discover the tail structure predicted by some models, future observations should cover more baselines."
    },
    "0206/astro-ph0206295_arXiv.txt": {
        "abstract": "Star clusters in 6 nearby spiral galaxies are examined using archive images from the Wide Field Planetary Camera 2 (WFPC2) on board the Hubble Space Telescope (HST).  The galaxies have previously been studied from the ground and some of them are known to possess rich populations of ``young massive clusters'' (YMCs). Comparison with the HST images indicates a success-rate of $\\sim75$\\% for the ground-based cluster detections, with typical contaminants being blends or loose groupings of several stars in crowded regions.  The luminosity functions (LFs) of cluster candidates identified on the HST images are analyzed and compared with existing data for the Milky Way and the LMC. The LFs are well approximated by power-laws of the form $dN(L)/dL \\propto L^\\alpha$, with slopes in the range $-2.4\\la\\alpha\\la -2.0$. The steeper slopes tend to be found among fits covering brighter magnitude intervals, although direct hints of a variation in the LF slope with magnitude are seen only at low significance in two galaxies.  The surface density of star clusters at a reference magnitude of $M_V=-8$, \\scl[-8], scales with the mean star formation rate per unit area, \\ssfr .  Assuming that the LF can be generally expressed as $dN(L)/dL = c \\, A \\, \\ssfr^\\gamma \\, L^\\alpha$, where $A$ is the galaxy area, $\\gamma\\sim 1.0 - 1.4$, $\\alpha=-2.4$ and the normalization constant $c$ is determined from the WFPC2 data analyzed here, the maximum cluster luminosity expected in a galaxy from random sampling of the LF is estimated as a function of \\ssfr\\ and $A$.  The predictions agree well with existing observations of galaxies spanning a wide range of \\ssfr\\ values, suggesting that sampling statistics play an important role in determining the maximum observed luminosities of young star clusters in galaxies. ",
        "introduction": "Since the launch of the Hubble Space Telescope (HST), the ubiquity of highly luminous young star clusters in starburst environments, sometimes referred to as ``super star clusters'' (SSCs), has been firmly established. Some of the best-known cases include merger galaxies such as the ``Antennae'' \\citep{ws95} and NGC~3256 \\citep{zepf99}, but SSCs have also been found in galaxies that do \\emph{not} show any obvious indications of having been involved in recent merger events.  Examples of isolated starburst dwarfs with SSCs are NGC~1569 and NGC~1705 \\citep{as85,oc94}. The Large Magellanic Cloud also contains a number of ``blue populous'' or ``young massive'' clusters (YMCs) \\citep{sn51,hod61}, many of which have ages characteristic of Milky Way \\emph{open} clusters, i.e.\\ a few times $10^8$ years or younger, but are an order of magnitude more massive than any young cluster that is known in our Galaxy today.  Such clusters are also present in some spiral galaxies, and there appears to be a fairly tight relation between the star formation rates (SFRs) of galaxies and the number of YMCs \\citep[][hereafter LR2000]{lr00}.  The current terminology is confusing and reflects, to a large extent, a lack of understanding of how different types of clusters are related.  The terms ``populous'', ``massive'' and ``super'' clusters were introduced as a reference to star clusters that apparently have no counterpart in our Galaxy, but do they actually constitute a separate class of objects, fundamentally different from low-mass ``open'' clusters? Do they require special conditions to form? There is not really any \\emph{a priori} reason to assume that the open cluster system in the Milky Way is representative for other galaxies or even other spirals, and by the same token, the \\emph{presence} of YMCs in e.g.\\ the LMC may be no more unusual than their \\emph{absence} in our Galaxy. It may be more appropriate to view the cluster system of our Galaxy as part of a continuum that ranges from very cluster-poor galaxies with low SFRs such as IC~1613 \\citep{wyd00}, over our Galaxy, to the LMC and finally to starburst environments.  Simple statistics may play a role in determining the luminosity of the brightest cluster in a galaxy, since random sampling of a power-law luminosity function (LF) will give brighter clusters in galaxies with richer cluster systems \\citep{whit01}. Rich cluster systems, in turn, would generally be expected in galaxies with high levels of star formation.  In addition to statistical effects there may be a physical upper limit to the mass of star clusters that can form in a galaxy, determined by factors such as the gas density and -pressure of the interstellar medium which are also expected to correlate with the overall star formation rate \\citep{ee97,ken98,bhe02}. It is currently unclear to what extent sampling statistics or physics dominate the maximum observed cluster luminosities and -masses, and under what conditions one or the other might prevail.  In reality, very little is known about how the LF of young cluster populations depends on environment.  \\citet{vl84} found that the LF of Milky Way open clusters is well described as a power-law of the form $dN(L)/dL \\propto L^\\alpha$ with $\\alpha=-1.5$ for $-8<M_V<-2$, but may steepen somewhere in the range $-11<M_V<-8$. Studies of young cluster populations in other galaxies generally find somewhat steeper slopes, $-2.5 \\la \\alpha \\la -2.0$ \\citep[e.g.][]{dk02,zepf99}, and \\citet{whit99} suggested a bend at $M_V\\sim-10.4$ for young clusters in the Antennae with $\\alpha=-2.6\\pm0.2$ for brighter clusters and $\\alpha=-1.7\\pm0.2$ for fainter clusters.  Unfortunately, there is little overlap between observations of open clusters in the Milky Way, which cover a very limited section of our Galaxy and include few clusters \\emph{brighter} than $M_V\\approx-9$, and studies of young clusters in external galaxies where the identification of clusters \\emph{fainter} than $M_V\\approx-8$ is difficult because of the risk of confusion with individual, luminous stars. Even with HST, unambiguous identification of star clusters in crowded environments is problematic at distances greater than 10--15 Mpc.  Considering the diversity of properties encountered among young cluster populations even in normal spirals, HST images of nearby spirals may provide some insight into correlations between properties of the cluster systems and host galaxy parameters. The HST archive contains a number of WFPC2 datasets covering nearby spiral galaxies, including many of those studied by LR2000. At the typical distances of these galaxies ($\\approx 5$ Mpc), one WF pixel corresponds to a linear scale of 2--3 pc, comparable to the typical half-light radius of a star cluster.  This makes it possible to identify star clusters based on their angular extent and thereby greatly reduce the risk of contamination by individual stars.  In this paper I examine HST archive images of 6 selected galaxies from LR2000, some of which possess rich populations of YMCs while others have more modest, presumably more Milky Way-like cluster populations. As with any survey based on archive data, the choices of exposure times, filters and exact pointing are limited by what happens to be available. However, even with exposures in only one band one can perform interesting comparisons of the cluster populations in different galaxies. As a bonus, the WFPC2 data can be used for a check of the ground-based cluster identifications. The detailed structure of individual clusters will be investigated in a forthcoming paper. ",
        "conclusions": "The cluster populations of 6 nearby spiral galaxies have been examined on archive HST/WFPC2 images.  Reidentification of cluster candidates selected on ground-based images in previous work suggests a success-rate of roughly 75\\% for ground-based cluster selection.  The ground-based photometry tends to overestimate the brightness of the clusters by up to several tenths of a magnitude, because other objects than the cluster itself are included within the relatively large apertures used for the ground-based photometry. On the other hand, ground-based colors agree very well with the WFPC2 photometry.  Typical slopes for the cluster luminosity functions $dN(L)/dL \\propto L^{\\alpha}$ studied here are $-2.4 \\la \\alpha \\la -2.0$, clearly steeper than the $\\alpha=-1.5$ for the Milky Way open cluster LF found by \\citet{vl84} but in good agreement with data for LMC clusters taken from \\citet{bica96}.  However, the Milky Way data cover a much fainter magnitude interval and van den Bergh \\& Lafontaine pointed out that the lack of very luminous clusters in the Milky Way by itself suggests a steepening of the LF at the bright end.  There is a general tendency for fits covering brighter luminosity intervals to have steeper slopes, but a direct hint of a change in slope at $M_V\\approx-8$ is seen only in two galaxies (NGC~1313 and NGC~5236), at the $\\la1\\sigma$ level.  The surface densities of clusters ($\\scl[-8]$) scale with the area-normalized star formation rates of the host galaxies.  Assuming that this scaling is generally valid, the maximum cluster luminosity expected from random sampling of the LF was estimated by extrapolation of a power-law LF with $\\alpha=-2.4$, scaled to the area $A$ and \\ssfr\\ for a sample of galaxies. Both the slope and scatter of the predicted relation agree well with observations by \\citet{lr00} and \\citet{bhe02}, spanning a wide range in \\ssfr\\ and physical dimensions.  The maximum cluster luminosities are, however, overpredicted by about 0.9 mag, possibly because the mean density of clusters in the galaxies is lower than estimated from the HST pointings.  Some starburst dwarfs do manage to form clusters that are significantly more luminous than expected from the size of their cluster populations, indicating that the LF of young clusters in these galaxies may be peculiar.  It remains to verified whether the cluster LF is really universal, concerning the slopes and their dependence on magnitude, as well as the location and even the existence of a ``bend'' in the LF. Data presented here suggest that a change in slope occurs at $M_V\\sim-8$, but observations of a similar bend at $M_V\\sim-10.4$ in the ``Antennae'' suggest a possible dependency on factors such as environment and/or the age distribution of the clusters involved.  An interesting point is whether or not the mass distribution observed in young cluster systems can be reconciled with that of old globular clusters (GCs). The low-mass end of the old GC mass spectrum is likely to be strongly affected by dynamical evolution \\citep{fz01,ves00,ves01}, but for masses above $\\sim10^5$ \\msun\\ the mass distribution is well approximated by a power-law with slope $\\sim-1.9$ \\citep{kis94,kis96,lar01b}.  This is somewhat shallower than the slopes of the luminosity functions at bright magnitudes found in this paper, but similar to the faint-end slopes.  The interpretation of these differences is complicated by the fact that the LF of a population of clusters with a range of ages is, in general, different from the more fundamental underlying \\emph{mass} distribution.  In particular, if the mass function has a physical upper limit then the LF of a cluster population with a range of ages will tend to have a steeper slope than the underlying mass distribution.  This is well illustrated by the finding that the mass function of Antennae clusters follows a single power-law with $\\alpha=-2$ over the range $10^4 < M < 10^6\\,\\msun$ \\citep{zf99}.  Because random sampling statistics appear to play an important role in determining the luminosity of the brightest clusters observed in galaxies, the role of physical processes in regulating the upper mass limit of star clusters in galaxies is difficult to determine. For galaxies with a continuous star formation history, observations of the cluster \\emph{mass} distribution should make it possible to establish with greater certainty how the maximum possible cluster mass depends on host galaxy properties. This, however, would require fairly accurate age estimates for large samples of clusters, down to $M_V\\sim-6$ or fainter. The new Advanced Camera for Surveys on board HST will be very suitable for such studies because of its large field of view and superior sampling compared to WFPC2.  An essential ingredient in future surveys should be inclusion of $U$ band imaging, which would allow better age estimates as well as providing constraints on cluster reddenings. Multi-color photometry would also make the cluster identifications more secure."
    },
    "0206/astro-ph0206226_arXiv.txt": {
        "abstract": "Using a redshift survey of \\ncz galaxies (\\ncznew new or remeasured) in a region of 95 square degrees centered on the nearby galaxy cluster Abell 2199, we analyze the supercluster containing A2199, A2197, and an X-ray group.  The caustic technique accurately reproduces the true mass profiles of simulated simple superclusters (i.e., superclusters where the virial mass of one cluster is 2-10 times the virial mass of all other clusters in the supercluster).  We calculate the masses of the two main components of A2197 (A2197W and A2197E) using archival X-ray observations and demonstrate that the A2199 supercluster is simple (the mass of A2199 is 5 and 12 times larger than A2197W and A2197E respectively) and thus that the caustic technique should yield an accurate mass profile.  The masses of A2199, A2197W, and A2197E within $r_{500}$ (the radius within which the enclosed density is 500 times the critical density) are 22.0, 3.8, and 1.7$\\times 10^{13}\\msun$.  The  mass profile is uncertain by $\\sim$30\\% within $3~\\Mpc$ and by a factor of two within $8~\\Mpc$ and is one of only a few for a supercluster on such large scales.  Independent X-ray mass estimates agree with our results at all radii where they overlap. The mass profile strongly disagrees with an isothermal sphere profile but agrees with profiles suggested by simulations.  We discuss the interplay of the supercluster dynamics and the dynamics of the bound subclusters.  The agreement between the infall mass profile and other techniques shows that the caustic technique is surprisingly robust for simple superclusters. ",
        "introduction": "Clusters of galaxies are important probes of the distribution of matter on intermediate scales.  Clusters are surrounded by infall regions where the galaxies are falling into the gravitational potential well of the cluster, but have not yet reached equilibrium. Many, perhaps most, of the galaxies in this region are on their first pass through the cluster. They populate a regime between that of relaxed cluster cores and the surrounding large-scale structure where the transition from linear to non-linear clustering occurs. Recent papers have explored infall regions using two-body dynamics of binary clusters \\citep[e.g., ][]{mw97}, the virial theorem in superclusters \\citep{small98}, and weak lensing \\citep{kaiserxx}. \\citet{ellingson01} demonstrate that these infall regions are the site of significant galaxy evolution and are key to understanding the Butcher-Oemler \\citep{1984ApJ...285..426B} effect.  Two recent papers \\citep{rqcm, rines01b} demonstrate that the infall regions of clusters (the Shapley Supercluster and Abell 2199 respectively) provide evidence of hierarchical structure formation on supercluster scales.   Because clusters are not in equilibrium outside the virial radius, neither X-ray observations nor virial analysis provide accurate mass determinations at large radii. There are now two methods of approaching this problem: weak gravitational lensing \\citep{kaiserxx} and kinematics of the infall region \\citep[][hereafter DG and D99 respectively]{dg97,diaferio1999}.  In redshift space, the infall regions of clusters form a characteristic trumpet-shaped pattern. These caustics arise because galaxies fall into the cluster as the cluster potential overwhelms the Hubble flow \\citep{kais87,rg89}. Under simple spherical infall, the galaxy phase space density becomes infinite at the caustics. DG analyzed the dynamics of infall regions with numerical simulations and found that in the outskirts of clusters, random motions due to substructure and non-radial motions make a substantial contribution to the amplitude of the caustics which delineate the infall regions \\citep[see also][]{vh98}. DG showed that the amplitude of the caustics is a measure of the escape velocity from the cluster; identification of the caustics therefore allows a determination of the mass profile of the cluster beyond the virial radius. DG find that nonparametric measurements of caustics in simulated clusters yield mass profiles accurate to $\\sim$50\\% on  scales up to $10~\\Mpc$.  The caustic technique places no requirements on the equilibrium state of the cluster; rather, it assumes only that galaxies trace the velocity field. Indeed, simulations suggest that little or no velocity bias exists on linear and mildly non-linear scales \\citep{kauffmann1999a,kauffmann1999b}.  \\citet{vh98} used simulations to explore an alternative parametric analysis of the infall region using a maximum likelihood method. Their technique requires assumptions about the functional forms of the density profile and the velocity dispersion profile.  \\citet[][hereafter GDK]{gdk99} applied the kinematic method of DG to the infall region of the Coma cluster. They successfully reproduced the X-ray derived mass profile and extended direct determinations of the mass profile to $\\sim$10$~\\Mpc$.  This method has also been applied to the Shapley Supercluster \\citep{rqcm}, A576 \\citep{rines2000}, A1644 \\citep{tustin} and the Fornax cluster \\citep{drink}.   We apply the method of DG and D99 to A2199, one of the richest, most regular nearby clusters of galaxies \\citep[][hereafter MVFS]{mvfs}.  \\citet{rines01b} presented a preliminary view of this system and demonstrated that it contains several infalling, bound subclusters.  This system provides an important test of the robustness of the caustic technique in simple superclusters (i.e., those dominated by a cluster 2-10 times more massive than any other system in the supercluster).   We describe numerical simulations of simple superclusters in $\\S$ 2 to demonstrate the validity of the caustic technique for these systems. We describe the spectroscopic observations in $\\S$ 3.  In $\\S$ 4, we show that A2197 is a superposition of groups and demonstrate that A2199/A2197 is indeed a simple supercluster.  We determine the amplitude of the caustics, calculate the mass profile, and compare the infall mass profile to other estimators in $\\S 5$. We discuss our results in $\\S 6$ and conclude in $\\S 7$.  The physical scale at the redshift of the supercluster ($cz=9156~\\kms$) is $1^\\circ = 1.54~\\Mpc$ ($H_0 = 100~h~\\kms, \\Omega _m = 0.3, \\Omega _\\Lambda = 0.7$). ",
        "conclusions": "We have previously demonstrated that the infall region of A2199 contains several bound subclusters \\citep{rines01b}.  Here, we show that, in simulations, caustics yield accurate mass profiles of superclusters dominated by a central cluster 2-6 times more massive than all other systems in the supercluster.  The mass profiles of the simulated clusters agree with NFW or Hernquist models but not with a singular isothermal sphere.  X-ray data demonstrate that A2197 is the superposition of two bound X-ray groups (centered on the bright galaxies NGC 6160 and NGC 6173) which may be interacting.  Because of the complexity of A2197, the virial theorem does not yield an accurate mass estimate.  The X-ray masses for A2197W and A2197E are factors of 6 and 13 smaller than A2199.  Tests of similar superclusters from simulations indicate that the caustic technique should yield an accurate mass profile for the A2199 supercluster.  The supercluster mass profile suffers some uncertainty due to the complexity of the surrounding large-scale structure and the resulting difficulty in determining the redshift center of the supercluster. Within $3~\\Mpc$, however, this uncertainty is no larger than the scatter predicted by numerical simulations for isolated clusters. The infall mass profile is in excellent agreement with independent X-ray estimates at small radii. At large radii, the mass profile is clearly inconsistent with an isothermal sphere but agrees with profiles motivated by simulations.  The mass profile indicates that A2197W, A2197E, and NRGs388 are bound, infalling groups; NRGs385 and NRGs396 (and the NGC 6159 group for inclination angles $\\gtrsim 20^\\circ$) are not infalling onto the supercluster in the present epoch.  This result demonstrates that the presence of caustics at large radii from clusters does not prove that galaxies or groups within these caustics are inside the infall region. In other words, the gravitational potential of a cluster and surrounding large-scale structure can significantly affect the kinematics of objects outside the turnaround radius.  The observed velocity dispersions of the groups surrounding A2199 are significantly larger than predicted by their X-ray properties and scaling relations.  The observed kinematic properties of galaxy groups within and near superclusters are likely dominated by the supercluster and not by the group potential.  The caustic technique is a robust estimator of cluster mass profiles ranging in mass from Fornax to the Shapley Supercluster.  We demonstrate that the technique is robust for simple superclusters, both in simulations and in observations of the A2199/A2197 supercluster.  We plan to analyze a sample of 8 nearby clusters and study the mass-to-light ratio throughout the infall regions of several of these systems."
    },
    "0206/astro-ph0206360_arXiv.txt": {
        "abstract": "We present the mass function (MF) of the Arches cluster obtained from ground-based adaptive optics data in comparison with results derived from HST/NICMOS data. A MF slope of $\\Gamma = -0.8 \\pm 0.2$ is obtained. Both datasets reveal a strong radial variation in the MF, with a flat slope in the cluster center, which increases with increasing radius. ",
        "introduction": "The Arches cluster is one of the most massive young star clusters found in the Milky Way. The total stellar mass is estimated to be $\\ga 10^4 M_\\odot$, with a central stellar density of $3 \\cdot 10^5 M_\\odot\\ {\\rm pc}^{-3}$ (Figer et al. 1999). Located at a projected distance of only 25 pc from the Galactic Center (GC), it evolves under extreme physical conditions. As the formation of massive stars ($M > 30 M_\\odot$) is proposed to be strongly dependent on high stellar and gas densities (Behrend \\& Maeder 2001, Bonnell et al. 1998), star clusters close to the GC are important test beds for star formation (SF) scenarios in dense regions. Displaying the highest stellar density found in a massive young cluster (YC) in the Milky Way, the Arches cluster is a unique object for the comparison with massive SF models. Due to the proximity to the GC, strong tidal forces act towards the disruption of the cluster entity, causing the timescale for dynamical evolution to be short, which may lead to the dissolution of the cluster within $10-20\\,{\\rm Myr}$ (Kim et al. 1999, Portegies Zwart et al. 2002). With an age of only $2.5 \\pm 0.5$ Myr (Figer et al. 2002, in prep.), the stellar population of the Arches cluster should be mostly unaffected by stellar evolution, but the spatial appearance of the cluster should reflect the fast dynamical evolution. ",
        "conclusions": "The radial variation in the MF of the Arches cluster as well as the flat integrated mass function support massive cluster and star formation models, suggesting high-mass stars to form in dense cluster environments. Although the radial variation in the MF strongly suggests mass segregation in Arches to be present, the short dynamical timescales, strongly influenced by the GC tidal field, prohibit discrimination between primordial and dynamical segregation effects.  The close similarity of the mass functions obtained from the Gemini AO and the HST datasets reveal that high-resolution ground-based AO data are capable to produce comparable physical results as space-based observations - in particular in dense regions where spatial resolution is essential."
    },
    "0206/astro-ph0206156_arXiv.txt": {
        "abstract": "{ We present several new sets of grids of model stellar atmospheres computed with modified versions of the ATLAS9 code. Each individual set consists of several grids of models with different metallicities ranging from \\logZ\\ $= -$2.0 to +1.0 dex. The grids range from 4000 to  10000~K in \\Teff\\ and from 2.0 to 5.0 dex in \\logg. The individual sets differ from each other and from previous ones essentially in the physics used for the treatment of the convective energy transport, in the higher vertical resolution of the atmospheres and in a finer grid in the (\\Teff, \\logg) plane. These improvements enable the computation of derivatives of color indices accurate enough for pulsation mode identification. In addition, we show that the chosen vertical resolution is necessary and sufficient for the purpose of stellar interior modelling. To explain the physical differences between the model grids we provide a description of the currently available modifications of ATLAS9 according to their treatment of convection. Our critical analysis of the dependence of the atmospheric structure and observable quantities on convection treatment, vertical resolution and metallicity reveals that spectroscopic and photometric observations are best represented when using an inefficient convection treatment. This conclusion holds whatever convection formulation investigated here is used, i.e.\\ MLT($\\alpha=0.5$), CM and CGM are equivalent. We also find that changing the convection treatment can lead to a change in the effective temperature estimated from Str\\\"omgren color indices from 200 to 400~K. ",
        "introduction": "Convective transport of energy in a stellar atmosphere is one of the most complex astrophysical problems. Many of the approximations usually admitted for the stellar interior, such as diffusive radiative transfer, are no longer valid. Moreover, throughout most of a convective stellar atmosphere radiative losses are large enough to make convection less efficient in transporting energy than radiation. Only stars which have a surface convection zone (CZ) extending deep into the stellar envelope can maintain efficient convective energy transfer near the bottom of their atmosphere. On the other hand, inefficient convection appears in all stars near the boundary of a convection zone close to locally stable regions. The modelling of inefficient convection requires a detailed knowledge about the effect of radiative gains and losses on the fluid flow. The situation is particularly complex for stars which are cool enough to develop a granulation pattern, such as the sun. In this case, at identical geometrical depths, vastly different physical conditions may be encountered depending on whether upflow in a granule or downflow in an intergranular lane is considered. The former may be optically thick while the latter is already optically thin, a consequence of the extreme temperature sensitivity of the dominant opacity source in the solar photosphere, the H$^{-}$ ion \\citep[cf.\\ also][]{Stein98}.  Currently, only very simple convection models are available for routine computation of extended grids of model atmospheres, while detailed numerical simulations are still unaffordable for applications that require the calculation of many thousands of individual model atmospheres over the HR diagram.  Our intention here is first to review the convection models which are available for use together with the popular ATLAS9 model atmosphere code by \\citet{kuruczI,kuruczII} \\citep[see also][]{cast97}. We provide an overview on what is known about the effects of the different convection treatments on model atmosphere structure and consequently on observable quantities.  The second purpose of the present paper is to determine to what extent the precision of fundamental parameters derived from the observed stellar spectrum, i.e.\\ \\Teff, gravity and metallicity depend on the model atmosphere.  Another objective is to obtain very accurate colors and more importantly very accurate derivatives of colors, color indices and limb darkening coefficients. These quantities are needed in the procedure of pulsation mode identification which is the first and a crucial step in any seismological study. Indeed probing the stellar interior of a pulsating star requires the knowledge of the resonant cavity within which each mode propagates, i.e.\\ the physical nature of the pulsation mode associated with each observed oscillation frequency. One such procedure is based on the computation of oscillation amplitude ratios and phase differences which in turn depend on the variation of the colors with effective temperature and gravity. The results of this application of the model atmosphere grids will be presented in the next papers of this series \\citep{Bar2002,Gar2002}. Finally, due to their enhanced resolution the new model grids are also useful to improve the outer boundary conditions of stellar structure calculations \\citep{Mont2001,DAntona2002}.  These goals are part of a program performed in the framework of preparing the COROT space mission (see \\citeauthor{COROT}). To achieve these purposes, we have used the ATLAS9 code in several versions modified for the convection zone treatment to compute new grids of model atmospheres, corresponding fluxes, surface intensities, $uvby$ colors, synthetic spectra for some representative lines, and compared them with relevant observations. We have three versions of the ATLAS9 code at our disposal: \\begin{enumerate} \\item The original version from CDROM13 of Kurucz \\citep{kuruczI} in which the convection zone is treated using mixing length theory (MLT). While ATLAS versions from 5 to 8 remained basically close to the formulations given in \\citet{BV58} and in \\citet{CG68}, some improvements were added in ATLAS9 \\citep[cf.][]{cast97}. In Sect.~\\ref{previous} we discuss the reasons for our specific selection among these improvements. \\item The other two versions were provided by one of the authors (FK) who modified the code to include turbulent convection models from \\citet[ CM]{cm}, and from \\citet*[ CGM]{cgm}. \\end{enumerate} Each convection model has been extensively used in the model atmosphere grid computations which we describe below. All the convection models are of local type and thus require the prescription of a characteristic length scale. Formally, it is possible to interchange the different length scales associated with the convection models. The motivation for doing so and a particular example will be discussed in the next paper of this series \\citep{Kup2002}.  This paper is organized as follows. In Sect.~\\ref{previous} we review previous works about the effect of the model structure on theoretical photometric colors and justify the need for new grids of model atmospheres. In Sect.~\\ref{convection} we describe the specific different convection treatments used and discuss their physical content.  In Sect.~\\ref{grids} we give details of the grid computations. In Sect.~\\ref{effects} we set out and comment the role of the convection treatments and convection parameters on the model structure, as well as its dependence on effective temperature, surface gravity, and metallicity. Finally, we discuss the consequences on observable quantities such as Balmer line profiles, flux distributions, and colors. ",
        "conclusions": "One of the main conclusions to be raised from this study is that as long as one considers inefficient convection, whatever is the choice of the formulation, either MLT with low $\\alpha$, or FST, the interpretation of spectroscopic or photometric observations is equivalent: observed BLPs and Str\\\"omgren color indices of dwarf and subgiant stars between A5 and G5 spectral types, and in a large range of metallicity are best represented by the use of less efficient convection transport, i.e.\\ MLT with $\\alpha = 0.5$, or with FST formulation. This confirms results already obtained by \\citet{fag93}, \\citetalias{cvm} and \\citet{Veer:98} for the Sun, Procyon, and other cool metal-poor stars using MLT models.  \\citet{Gard99} reported a few opposite cases (see Sect.~\\ref{previous}), but for parts of their sample of stars fundamentally known \\logg\\ values were not available. An analysis of a larger sample of stars in binary systems with revised fundamental parameters for {\\em both} \\logg\\ and \\Teff\\ \\citep{Smal:02} did not confirm the discrepancies previously found. Furthermore, for the case of F stars \\citet{Smal:02} noticed a larger systematic difference between fundamental effective temperatures and those obtained from H$_{\\beta}$ lines for MLT($\\alpha$=1.25) than for less efficient convection models, although this discrepancy remains within the overall uncertainties.  Nevertheless, we have to emphasize that in models with deep convection zones (e.g.\\ for Sun, Procyon) MLT($\\alpha = 0.5$) and FST treatments have comparable effects on calculated fluxes, but not on atmosphere structure. They produce different temperature gradients in the deep layers, as can be seen in Figs.~\\ref{Ttau_single} and \\ref{Ttau_general}, but those cannot be distinguished by the computed BLPs. In other words, the BLPs allow to discriminate among different values for the MLT parameter $\\alpha$, but not among MLT($\\alpha = 0.5$), CM, and CGM models. In any case, the sensitivity of BLPs to convection parameters depends significantly on the other physical parameters. This holds especially for the sensitivity to gravity change, which can be more important than usually expected.  In case of weakly efficient convection, fluxes and colors depend only weakly on the selected convection treatment. On the other hand, when the convection is highly efficient, then fluxes and colors become strongly dependent on the convection modelling, as the differences among the models show up more clearly within the photosphere. Thus, significant uncertainties on stellar global parameters arise from the convective treatment in model atmospheres. Ignoring these uncertainties can lead to systematic differences affecting subsequent interpretations.  The calculations of color and limb darkening partial derivatives are significantly improved when using the present model atmosphere grids which are finer spaced in \\Teff\\ and \\logg\\ and have a higher resolution in the temperature distribution with depth \\citep{Bar2002}. Smoothness of these derivatives is of crucial importance in the mode discrimination problem for non-radially pulsating stars, which is basically due to the dependence of the color amplitude ratios on these derivatives. Details of the required precision of these partial derivatives in order to be useful for mode identification will be given in \\citet{Gar2002}. There we will show that the next space asteroseismology missions -- COROT, MONS/R\\o mer and Eddington -- will supply light curves with high enough precision to permit a direct comparison up to the second order to partial derivatives with respect to temperature and gravity as calculated with the present model atmospheres.  The improved resolution of the new model grids also avoids unphysical oscillations in evolutionary track calculations when using ATLAS9 model atmospheres as boundary conditions (see Sect.~\\ref{resolution}). Moreover, the possibility to choose among different convection models allows a self-consistent match between model atmospheres and model envelopes \\citep{Mont2001}. However, we must stress here that the different relations $T$ and $F_{\\rm conv}/F_{\\rm tot}$ vs.\\ depth represent stars which are different in their radii and luminosities. The broad effects of the convective treatment can only be assessed by studying a complete stellar model, i.e.\\ a model with an atmosphere and an internal structure which are consistently built with the same convection formulation. We will address this topic in follow up work \\citep{Kup2002}."
    },
    "0206/astro-ph0206374_arXiv.txt": {
        "abstract": "We present results from a 47~ks observation of the Andromeda galaxy, M31, using the High-Resolution Camera of the Chandra X-Ray Observatory.  We detect 142 point sources spanning three orders of magnitude in luminosity, from $L_X = 2\\times 10^{35} \\rm \\, erg \\, s^{-1}$ to $L_X = 2\\times 10^{38} \\rm \\, erg \\, s^{-1}$ in the 0.1-10~keV band.  The X-ray source location accuracy is better than $1\\arcsec$ in the central regions of the galaxy.  One source lies within $1.3\\arcsec$ of SN 1885 but does not coincide with the UV absorption feature identified as the supernova remnant.  However, there is an optical transient, which is likely an optical nova, at the location of the X-ray source.   There is a weak source, $L_X \\sim 4 \\times 10^{36} \\rm \\, erg \\, cm^{2} \\, s^{-1}$, coincident with the nucleus of M31, and 14 sources coincident with globular clusters.  Our observation has very high efficiency down to luminosities of $1.5\\times 10^{36} \\rm \\, erg \\, s^{-1}$ for sources within $5\\arcmin$ of the nucleus.   Comparing with a ROSAT observation made 11 years earlier, we find that $0.46 \\pm 0.26$ of the sources with $L_X > 5 \\times 10^{36} \\rm \\, erg \\, s^{-1}$ are variable.  We find no evidence for X-ray pulsars in this region, indicating that the population is likely dominated by low-mass X-ray binaries.  The source density radial profile follows a powerlaw distribution with an exponent of $1.25 \\pm 0.10$ and is inconsistent with the optical surface brightness profile. The x-ray point source luminosity function is well fitted by a differential broken powerlaw with a break at a luminosity of $(4.5^{+1.1}_{-2.2}) \\times 10^{37} \\rm \\, erg \\, s^{-1}$.  The luminosity function is consistent with a model of an aging population of X-ray binaries. ",
        "introduction": "Study of X-ray sources in external galaxies is important for understanding the formation history and population statistics of X-ray binaries and other X-ray sources both in external galaxies and in our own.  Such studies will help us understand the evolutionary history of X-ray binaries, should provide information on the star formation history (e.g.\\ White \\& Ghosh 1998), and may be important in estimating the rate of merging objects (e.g.\\ Bethe \\& Brown 1999) critical for determining the rate of gravitational wave events.  The nearby, bright spiral galaxy M31 offers an excellent site for such studies.  The distance of the galaxy is well known, making luminosity estimation from flux measurement straight forward, the galaxy is sufficiently inclined so that sources can be reliably located within its morphology, and it is relatively nearby (780~kpc) so that many sources can be detected.  The sub-arcsecond resolution of the Chandra X-ray Observatory (CXO; Weisskopf 1988) permits individual X-ray sources to be discerned even in crowded regions of the galaxy and enables measurement of highly accurate positions.  Such position information is critical in finding unique optical and radio counterparts to the sources.  Here, we report on a deep observation of the core of M31 made with the Chandra High Resolution Camera (HRC; Murray et al.\\ 1997).  The HRC offers the best spatial resolution available for the X-ray study of M31.  We describe the observation and our analysis in \\S 2.  We discuss the variability in \\S 3 and  source identifications including a source very near, but not coincident with SN 1885, in \\S 4.  We consider the spatial and luminosity distributions of the population of X-ray point sources in \\S 5 and describe a model which relates the luminosity distribution to the age of the X-ray binary population.  Finally, we conclude in \\S 6 with comments on the implications of our results for understanding the formation history of X-ray binaries.   \\begin{figure*}[t] \\centerline{\\epsscale{1.6} \\plotone{f1_small.eps}} \\caption{Optical image of M31 from the Digital Sky Survey with positions of Chandra sources superimposed.  The sources are color-coded according to their luminosity: red for $L_X \\le 2 \\times 10^{36} \\rm \\, erg \\, s^{-1}$, magenta for $2 \\times 10^{36} \\rm \\, erg \\, s^{-1} < L_X \\le 2 \\times 10^{37} \\rm \\, erg \\, s^{-1}$, and blue for $L_X > 2 \\times 10^{36} \\rm \\, erg \\, s^{-1}$.  The green circle indicates the sources used in the analysis of group properties.  It is centered on the nucleus and and has a radius of $5\\arcmin$.  The arrow is $2\\arcmin$ long and points north. \\label{bigpic}} \\end{figure*}  \\begin{figure}[t] \\centerline{\\epsscale{1.0} \\plotone{f2.eps}} \\caption{Luminosity versus radial distance from the nucleus for X-ray point sources in M31. \\label{lum_r}}   \\end{figure} ",
        "conclusions": "Previous X-ray studies of M31 with pre-Chandra instruments indicated a good correspondence between the X-ray and optical profiles \\citep{tf91}.  However, the good correspondence does not continue to the small radii accessible with Chandra.  Instead, we find a powerlaw profile with an exponent of $1.25 \\pm 0.10$.  Powerlaw profiles with exponents near 1 are well known indicators of core collapse in globular clusters \\citep{djorgovski86}.  The same should be true for other stellar clusters such as the bulge of M31 \\citep{quinlan96}.  If the X-ray sources are predominately LMXBs which were formed in globular clusters \\citep{white02}, then the distribution may be a remnant of the inward migration and disruption of globular clusters \\citep{tremaine75}.  The X-ray point source luminosity function (XLF) of the core of M31 shows a distinct break near $4.4 \\times 10^{37} \\rm \\, erg \\, s^{-1}$. A break near this luminosity has been reported previously for M31 \\citep{primini93,shirey01}, although our analysis shows that the break luminosity found from fitting a broken powerlaw to the differential distribution is roughly a factor of two higher than that found from fitting to the cumulative distribution.  In the context of the model presented above \\citep{kilgard02}, the break luminosity is related to the age of the X-ray point source population.  Specifically, the X-ray population age $t_B = 1.3 \\eta \\bar{M}_2 /\\epsilon M_{\\odot} \\rm \\, Gyr$.  For accreting neutron stars in low-mass X-ray binaries, the efficiency $\\eta \\sim 0.1$ and the average companion mass $\\bar{M_2} \\sim M_{\\odot}$.  Good observational constraints on the duty cycle of emission are not available.  The duty cycles inferred for LMXBs in our Galaxy range from 1.0 to less than 0.01 \\citep{wijnands01}.  It has also been suggested, based on binary evolution models, that a population of LXMBs with low duty cycles is present in elliptical galaxies \\citep{piro02}.  The stellar population near the nucleus of M31 is one of the reddest known, leading to age estimates of 10~Gyr or older \\citep{rich95}.  If the X-ray population is of the same age, then $t_B \\gtrsim 10 \\rm \\, Gyr$ and we would estimate $\\epsilon \\sim 0.01$.  If the duty cycle were larger then the compact objects would accrete a total mass exceeding the companion mass in less than 10~Gyr."
    },
    "0206/astro-ph0206232_arXiv.txt": {
        "abstract": "Precise Doppler searches for extrasolar planets find a surfeit of planets with orbital periods of 3--4 days, and no planets with orbital periods less than 3 days. The circumstellar distance, $R_0$, where small grains in a protoplanetary disk reach sublimation temperature ($\\sim 1500$~K) corresponds to a period of $\\sim 6$ days.  Interior to $R_0$, turbulent accretion due to magneto-rotational instability may evacuate the disk center.  We suggest that planets with orbital periods of 3--4 days are so common because migrating planets halt once this evacuated region contains the sites of their exterior 2:1 Lindblad resonances. ",
        "introduction": "Analytic calculations \\citep{gold80,ward97a} and numerical simulations \\citep{nels00, kley01} suggest that protoplanets in a protoplanetary disk migrate rapidly into the star they orbit---so rapidly that it is a wonder any planets survive at all. Small protoplanets torque the disk at Lindblad and corotation resonances, and the resulting back-torque can propel a planet into the star in a matter of $10^{5}(M_P/M_{\\bigoplus})^{-1}$ years \\citep{ward97b}. \\citet{ward97b} has dubbed this conundrum the Shiva problem, after the Hindu god of destruction. Large protoplanets may open a gap in the disk via their resonant torques, and so become locked to the disk's viscous spreading, a process which may dump the most of the disk and planet onto the star within $10^{7}$ years or less, depending on the disk viscosity.  Figure~\\ref{fig:ahistogram} shows the distribution of the orbital periods of the innermost Doppler planet candidates, summarizing data from \\citet{exoplanets.org}\\footnote{This website about extrasolar planets is available at \\url{http://exoplanets.org/}}.  These candidate planets all orbit stars with masses in the range 0.7--1.4 $M_{\\odot}$. Figure~\\ref{fig:ahistogram} suggests that whatever mechanism halted the migration of these planets operates best at an orbital period of $\\sim 3$~days and ceases to operate at shorter periods.  Of the 20 planets with periods less than 20 days, 8 have periods in the range 3--4 days. No planet has a period less than 2.98 days. This trend appears to be real and not an artifact of observational selection;  the primary precise-doppler surveys are complete for Jupiter mass planets out to a period of $\\gtrsim 0.5$ years \\citep{butl01}.  \\begin{figure} \\epsscale{1.0} \\plotone{f1.ps} \\caption{Histogram of the orbital periods of the extrasolar planet candidates detected by the Precision-Doppler technique.} \\label{fig:ahistogram} \\end{figure}   Occasionally the interactions among two planets and a star can leave a planet trapped by stellar tides into a circular orbit at $\\sim 0.04$ AU \\citep{rasi96}.  Sometimes an accreting planet may overflow its Roch lobe, losing mass to the star, and this process may halt the planet's migration \\citep{tril98}. But these phenomena appear to be rare.  Scenarios in which planets migrate by interacting with a disk of planetesimals without gas \\citep{murr98} provide a disk truncation radius where planets may gather: the radius where planetesimals become hot enough to sublimate. The problem with gas-less migration schemes is that substantially changing the orbit of a Jupiter-mass planet requires roughly a Jupiter mass of planetesimals.  In contrast, optically-thick disks with more than a Jupiter mass of gas appear to be ubiquitous around young stars. \\citet{lin96} have suggested that planet migration ends where the gaseous protoplanetary disk meets the stellar magnetosphere.  We suggest an alternative explanation for the pile-up of planets near the 3-day period: a gas disk truncated at a temperature of 1500~K by the onset of Magneto-Rotational instability (MRI; Chandrasekhar 1961; Balbus \\& Hawley 1991). ",
        "conclusions": "We have painted a picture of the central regions of protoplanetary disks in which a low-surface-density, rapidly accreting central region catches inwardly-migrating planets like flypaper. A good way to test our conjecture that disk temperature---not the stellar magnetic field---determines the orbital radii of the innermost surviving planets is to extend the sample of planet search target stars to include a wider range of stellar masses. For example, we predict that planets around early A main sequence stars will collect at a radius much farther from the star ($\\sim 0.3$ AU) than planets around solar type stars.   The high rotation rates of main sequence A stars betray that magnetic disk truncation can not rescue planets migrating into these stars until they reach near 0.03 AU.  Doppler planet search techniques may not work for A stars because of their rotationally-broadened spectral lines. However, astrometric searches, transit searches, and direct-imaging searches for extrasolar planets are not so limited."
    },
    "0206/astro-ph0206248_arXiv.txt": {
        "abstract": "{We report the discovery of a strong correlation between the shape of a bulge's light-profile and the mass of its central supermassive black hole ($M_{\\rm bh}$).   We find that $\\log (M_{\\rm bh}/M_{\\sun}) = 2.91(\\pm0.38)\\log(n) +  6.37(\\pm0.21)$, where $n$ is the S\\'ersic $r^{1/n}$ shape index. This correlation is shown to be at least as strong as the relationship between the logarithm of the stellar velocity dispersion and $\\log M_{\\rm bh}$ and has comparable scatter.}  \\resumen{Presentamos el descubrimiento de una estrecha correlaci\\'on entre la forma del perfil de brillo de los bulbos y la masa del agujero negro supermasivo ($M_{\\rm bh}$) que contienen en su centro. Encontramos que $\\log (M_{\\rm bh}/M_{\\sun}) = 2.91(\\pm0.38)\\log(n) +  6.37(\\pm0.21)$, donde $n$ es el par\\'ametro de forma de la ley de S\\'ersic $r^{1/n}$. Esta correlaci\\'on es al menos tan fuerte como la relaci\\'on encontrada entre el logaritmo de la dispersi\\'on de velocidades estelares y $\\log M_{\\rm bh}$, y tiene una dispersi\\'on comparable.}    \\addkeyword{galaxies: fundamental parameters} \\addkeyword{galaxies: kinematics and dynamics} \\addkeyword{galaxies: nuclei} \\addkeyword{galaxies: structure}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206281_arXiv.txt": {
        "abstract": "We present  in this  paper the numerical  treatment of  the coupling between hydrodynamics and radiative  transfer.  The fluid is modeled by classical  conservation laws (mass, momentum and  energy) and the radiation  by the grey  moment $M_1$  system \\cite{m1}.   The scheme introduced  is able to  compute accurate  numerical solution  over a broad class of  regimes from the transport to  the diffusive limits. We propose an asymptotic  preserving modification of the HLLE scheme in order to treat  correctly the diffusion limit.  Several numerical results are presented,  which show that this approach  is robust and have the  correct behavior in both the  diffusive and free-streaming limits.  In  the last numerical example  we test this  approach on a complex physical  case by  considering the collapse  of a  gas cloud leading to  a proto-stellar  structure which, among  other features, exhibits very steep opacity gradients. ",
        "introduction": "Radiation hydrodynamics plays an important role in astrophysics, laser fusion and plasma physics. For many  years, efforts have been underway to develop mathematical  models   and  numerical schemes to     obtain accurate predictions at reasonable computing cost in this domain. \\\\  One  of   the  main  difficulty   is  to  obtain   accurate  numerical computations in  the various  regimes that can  be encountered  due to values  of material  opacities which  can vary  from several  order of magnitude. This can  be achieved by using the  full radiative transfer equation  but it  is still  out of  range for  complex  simulations in particular  in two  or  three spatial  dimensions.   To overcome  this difficulty several models  have been derived. For large  values of the material opacity, an asymptotic analysis leads to hyperbolic/parabolic systems of  equation refered to as the  equilibrium-diffusion limit or non  equilibrium  diffusion approximation  (\\cite{pom},\\cite{mihalas}, \\cite{lmh}, \\cite{dw1}). On the other  hand, for smaller values of the opacity  the streaming  limit can  be handled  by using  an hyperbolic system of equations coupling some  closure of the moment system of the transfer   equation   with   the  hydrodynamic   system   (\\cite{pom}, \\cite{mihalas},  \\cite{lmh},\\cite{bal})  .  Unfortunately for  several real applications,  such as radiative  shocks or star  formation, very different regimes  are encountered  in the same  simulation.  Coupling different models  on various zones  introduces large drawbacks  due to the domain partition and some loss of accuracy in the transitional zone.\\\\  From a numerical point  of view high-resolution schemes  are necessary to  take    into  account  the different    space   and  time  scales. Godunov-type schemes based on exact or approximate Riemann solvers are efficient   for  shock     problems  (\\cite{godlewski},   \\cite{toro}, \\cite{roe}).   In   radiation  hydrodynamics, without    source terms, hydrodynamic and radiation equations decouple so that a Riemann solver for the entire system can  easily be derived  from Riemann solvers for the     hydrodynamic       subsystem     and    for     the  radiation subsystem. Nevertheless, the difficulty  of this approach is to obtain the correct diffusion limit  for  the scheme  and correct behavior  in transitional regime.   Let us notice  that the regime at the numerical level depends  on the ratio of the  mean free path  of photons and the size of  the local cell.   An other difficulty is  related to the time scales in  the considered problems. In most  cases radiation has to be treated implicitly since the time step given by usual CFL condition is much  too small.  When   possible   the hydrodynamics may  be  treated explicitly (\\cite{dw2}), but  in some problems  the time step given by the  CFL  condition  on  the   hydrodynamic  subsystem  is still   too restrictive and hence a fully implicit approach is necessary. \\\\  In this  paper we propose new  model and numerical  scheme to describe radiation hydrodynamics in a wide range of regimes.  The equations for the  fluid are  the classical  conservation laws  (mass,  momentum and energy). The radiation  is described by the grey  moment system of the transfer equation (radiative energy and radiative flux) with the $M_1$ closure introduced in (\\cite{m1}).  This closure, based on the minimum entropy  principle,  is  able  to  describe large  anisotropy  of  the radiation  while  giving  also   the  correct  diffusion  limit.   The Eddington factor  is a  non constant function  of the reduce  flux and therefore the  radiation subsystem  is a nonlinear  hyperbolic system. Here, these  equations are written in the  comoving frame, introducing some non  conservative coupling terms. A Riemann  solver is developped by  performing  a wave  decomposition  neglecting the  nonconservative coupling terms  and the source terms,  which is trivial  since the two subsystems  then decouple.  Afterwards  the nonconservative  terms are reintroduced in  the scheme, by  using in the radiation  equations the velocity given  by the hydrodynamic solver. This  approach, natural in the streaming regime, can be  improved to capture the diffusion limit. For  that  purpose the  Riemann  solver  for  the $M_1$  subsystem  is modified  to take  into account  the stiff  relaxation term  along the ideas  developped  in   (\\cite{sjin}).   With  this  modification  our approach is  able to correctly treat  the streaming regime  as well as the  diffusion regime  with  the same  solver.   Finally as  radiative shocks  or star formation  simulations need  high resolution  and very different time scales we  describe the implementation of our radiation hydrodynamics solver in a moving grid-time implicit framework.  The outline  of this paper is  as follows.  In the  second section the equations of radiation hydrodynamics with the $M_1$-closure are given. In  section 3  we  present  the numerical  methods  for solving  these equations.   The  actual implementation  is  detailed  in  section 4. Numerical  experiments are  shown  in section  5,  first on  radiative shocks to  demonstrate the validity  of our method in  various regime, then on the collapse of a gas cloud leading to a stellar-like structure. The last section is the conclusion of the paper.\\\\ ",
        "conclusions": "We have presented  in this paper  a new numerical model for  radiative transfer which gives satisfactory results in  a wide range of physical conditions including both the diffusion and the free-streaming regimes where the   anisotropy  of the  photons distribution  function  can be large.  This method  is therefore well  suited  to modeled some  laser experiments   and many astrophysical flows.  Works  are in progress to incorporate  in the model a more  detailed description of the physical processes  (multigroup  approach,   ionization,  ....)  and to further develop this approach in a multidimensionnal framework."
    },
    "0206/astro-ph0206138_arXiv.txt": {
        "abstract": "{ Primordial superheavy particles, considered as the source of Ultra High Energy Cosmic Rays (UHECR) and produced in local processes in the early Universe, should bear some strictly or approximately conserved charge to be sufficiently stable to survive to the present time. Charge conservation makes them to be produced in pairs, and the estimated separation of particle and antiparticle in such pair is shown to be in some cases much smaller than the average separation determined by the averaged number density of considered particles. If the new U(1) charge is the source of a long range field similar to electromagnetic field, the particle and antiparticle, possessing that charge, can form primordial bound system with annihilation timescale, which can satisfy the conditions, assumed for this type of UHECR sources. These conditions severely constrain the possible properties of considered particles. }  \\PACS{*}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206412_arXiv.txt": {
        "abstract": "The XMM-NEWTON Science Archive (XSA) is being developed within the XMM-NEWTON Science Operations Centre at the ESA VILSPA Satellite Tracking Station of Villafranca del Castillo, near Madrid in Spain. The first public release is planned for April 2002.  Based on the ISO Data Archive architecture and design, the XSA will be accessible from the World Wide Web through a Java interface at http://xmm.vilspa.esa.es/xsa that will allow: \\begin{itemize} \\item powerful and complex queries against the observations and exposure catalogue for the first release and against the source catalogue in future releases \\item configurable results display, including product visualization tools \\item customisable product retrieval via a shopping basket \\item selection of product level \\item product retrieval via FTP or CDROM \\item proprietary rights protection for both display and data request \\end{itemize}  The XSA will also be ready to provide inter-operability with other archives and applications. ",
        "introduction": "The XMM-NEWTON data products are currently delivered by mean of CDROM to the observation Principle Investigators (PIs) using internal systems of the Science Operations Centre in VILSPA.  With the first proprietary data becoming public soon, there arose the need of a more general and user friendly data access and distribution system.  In early 2001, it was decided to build a new system called the XMM-NEWTON Science Archive (XSA) to allow general astronomers easy access to available XMM-NEWTON data products.  A set of representative users from the SOC, the SSC (Survey Science Centre, who are responsible for the generation of XMM-NEWTON data products) and some external X-ray scientists met in March 2001 to define the user requirements for the XSA. Matteo Guainazzi was nominated XSA Archive Scientist. His responsibility is to gather all users' feedback and requests and to consolidate them into the User Requirement Document (URD). He defines the implementation priorities for the software development team. Interacting daily with the developers, he makes sure that the XSA is developed according to the users' expectations. The first issue of XSA URD was released in early April 2001 for the design and development of the archive to start in late April 2001. The goal was to have a public release ready as soon as the first proprietary data becomes public, i.e. April 2002.  \\begin{figure}[ht] \\begin{center} {\\tt see Fig1.gif} \\end{center} \\caption{ XSA Query Panel } \\label{carviset-WA3_fig:fig1} \\end{figure}  In order to be able to meet such a short deadline, it was decided to adopt the open and flexible JAVA-based architecture of the successful ISO Data Archive (IDA, available at  http://www.iso.vilspa.esa.es/ida) which would offer a fast and cost-efficient development while ensuring that all X-ray community specific user requirements are fulfilled. ",
        "conclusions": "The XMM-NEWTON Science Archive developed at ESA provides a new, easy and fast access to all available XMM-NEWTON data products and auxiliary data through a state-of-the art powerful Java interface available from April 2002 at http://xmm.vilspa.esa.esa/xsa. It allows all the X-ray astronomers to have quick access to the data and exploit the science observed by XMM-NEWTON.  After the fully functional first public release in April 2002, the XSA will continue to be improved to offer more capabilities and access to newer data products. Work will also be done to ensure the XSA's full integration into the Virtual Observatory projects."
    },
    "0206/hep-ph0206257_arXiv.txt": {
        "abstract": "\\noindent It has recently been suggested that the existence of bare strange stars is incompatible with low scale gravity scenarios.  It has been claimed that in such models, high energy neutrinos incident on the surface of a bare strange star would lead to catastrophic black hole growth.  We point out that for the flat large extra dimensional case, the parts of parameter space which give rise to such growth are ruled out by other methods.  We then go on to show in detail how black holes evolve in the the Randall-Sundrum two brane scenario where the extra dimensions are curved.  We find that catastrophic black hole growth does not occur in this situation either.  We also present some general expressions for the growth of five dimensional black holes in dense media. ",
        "introduction": "The idea that the geometry of extra dimensions might be responsible for the hierarchy between the scale of electroweak physics and the Planck scale is extremely interesting.  In these models, the mass scale associated with gravity is around a TeV but appears to be much higher due to the small overlap of the extra dimensional graviton wave function with our standard model brane \\cite{add,randall}. In a gravity theory with $4+d$ space-time dimensions and a fundamental scale $M_F$, one expects black hole production at energy densities higher than $M_F^{4+d}$, so there has been a great deal of interest in the possibility of black hole production at the next generation of super-TeV scale colliders \\cite{dimopoulos,cheung,Giddings:2002bu}. The idea that colliders might produce small black holes is at first alarming, but these black holes are so small that they are expected to evaporate via Hawking evaporation before they are able to interact with their surroundings and grow.  A different situation would arise if the black hole were produced in an extremely dense medium like the interior of a neutron star, as in that case the black hole might interact with another particle before it decays, so that the Hawking evaporation would be balanced by the accretion of matter and the black hole might start to grow.   Production of the initial black hole requires that a nucleon belonging to the star be hit by an incident highly energetic particle such as a cosmic ray or a cosmic neutrino, with an energy of at least a few PeV to reach the threshold of black hole production, $\\sqrt{2\\ m_N \\ E_i} \\sim M_{BH} \\sim$ few TeV.  According to the hoop conjecture, the cross section for black hole production can be taken to be $\\sigma_{BH} = \\pi {r_s}^2$ where $r_s$ is the Schwarzschild radius of the centre of mass energy of the incident particle and the target.  Cosmic neutrinos could be a candidate for black hole production since $\\sigma_{BH}$ dominates over all the Standard Model neutrino-nucleon interactions for neutrino energies above $\\sim100$ PeV \\cite{Feng:2001ib}. Ultra High Energy neutrinos are expected to exist (as well as the already observed UHE cosmic rays \\cite{AGASA,FLYEYE}), although the current sensitivity of neutrino telescopes does not enable us to detect them \\cite{nutelescopes}. The most straightforward mechanisms of production would be via the interaction of UHE cosmic rays with the cosmological microwave background (GZK mechanism \\cite{GZK}) and via collisions of accelerated hadrons and photons inside astrophysical objects such as Active Galactic Nuclei. Other, more exotic production processes involving ``hidden sources'' or decay of ultra-heavy relic particles have also been proposed, possibly giving rise to many neutrinos with energies as high as $10^{22-23}$ eV \\cite{fluxes}. We prefer however to retain a more conservative estimate of the high energy neutrino flux, essentially based on the assumption that neutrinos are produced by the same cosmologically distributed extra-galactic sources that would be responsible for the observed high energy cosmic rays: the Waxman-Bahcall bound \\cite{WB}.  Using this bound, one can deduce the number of neutrinos of energy $E_{min}<E<E_{max}$ falling on a star of radius $R$ per unit time \\begin{equation} \\dot{N} = 800 \\ \\pi^2 \\left ( \\frac{R}{1 \\ km}\\right )^2 \\ \\left (\\frac{1\\ {\\rm GeV}}{E_{min}} - \\frac{1\\  {\\rm GeV}}{E_{max}}\\right )\\ \\ \\   s^{-1}. \\label{waxman} \\end{equation} This rate would become comparable to the corresponding expression for cosmic rays as the energy increases.  For the surface of a star with a radius of 10 km this rate is of about 40 neutrinos per year with an energy between $10^{20}$ eV and $2\\times 10^{20}$ eV, while the current measurements made on Earth, although still quite imprecise, would imply approximately 5 to 20 cosmic rays per year around $10^{20}$ eV.  A recent paper \\cite{learned} has pointed out that such a black hole formed by high energy neutrinos on the outside of a neutron star will not in fact grow since the region in which the black hole first forms is not dense enough for the black holes to interact with more nucleons before it decays.  The same paper also shows that the situation is fundamentally different in the case of strange stars.  It is postulated that the energy per quark in normal (up and down) quark matter may be higher than that in strange quark matter \\cite{witten} so it has been hypothesised that one possible end point of stellar evolution is a star entirely composed of up, down and strange quarks \\cite{olinto}.  Since it is thought that some strange stars may be `bare' in as much as their density rises from zero to more than nuclear densities in a length scale of order $\\sim$ fm \\cite{olinto}, strange stars could provide a medium in which TeV scale black holes could be created and then subsequently grow.  One would not expect a strange star to possess spectral lines and also for it to have a different mass-radius relation and cooling rate to a neutron star.   There has recently been a lot of interest in a possible strange star candidate \\cite{drake} although it is not clear yet if the identification of this source as a strange star is correct \\cite{prasanna}.  The authors of \\cite{learned} go on to say that the existence of a bare strange star would place constraints on the number and size of extra dimensions.  This is because for large enough extra dimensions and conservative estimates of high energy neutrino fluxes, the growth of a neutrino-nucleon interaction induced black hole will continue until it consumes the whole of the star.  The constraints obtained in that paper are mainly for the case of $d\\le 2$ flat extra dimensions.  There are many other stronger constraints on the case of 2 extra flat dimensions from astrophysics and cosmology \\cite{2con}.  For the case of a single extra dimension which solves the hierarchy problem, it is not possible for the extra dimension to be flat, since this would require physics to be effectively five dimensional at lengths up to solar system scales.  One therefore requires a warped extra dimension as in the model of Randall and Sundrum.  Phenomenologically such theories are difficult to constrain since the graviton Kaluza-Klein mode masses are of order of the fundamental scale $\\sim$ TeV \\cite{randall}.  This is fundamentally different from the flat extra dimension scenarios where the Kaluza-Klein masses are far below a TeV and can therefore be excited at astrophysical energies.  Given the recent possible detection of strange star candidates, it is interesting to find out if the existence of strange stars would place any constraints upon such scenarios.  In this work we briefly review the 5D Randall-Sundrum model and the black holes that can be formed in this theory.  We then write down some general equations describing the evolution of 5D black holes in dense media.  We show why TeV scale black holes created at the surface of neutron stars do not continue to grow and then describe the growth of black holes in strange stars assuming the existence of a single warped extra dimension. ",
        "conclusions": "In this work we investigated the possibility that the Randall-Sundrum 2-brane model of TeV scale gravity is incompatible with the existence of bare strange stars due to the growth of black holes seeded by high energy neutrinos.  We pointed out that the growth of such black holes is suppressed in these models when they reach the radius corresponding to the curvature radius of the $AdS_5$ in between the branes.  We performed detailed simulations to see if black holes would grow and engulf the star.  We saturated the Waxman-Bahcall bound to find the highest energy neutrino that one could expect to have hit the oldest known strange star candidate \\cite{walter} in the time since it has been discovered.  In doing so, we found that the regions of parameter space where growth of such holes may be permitted has already been ruled out by accelerator experiments.  In the paper of \\cite{learned} it was shown that catastrophic black hole growth in strange stars could only work if there were 1 or 2 large extra dimensions.  The case of 2 large flat extra dimensions has already been tightly constrained by astrophysical and cosmological constraints \\cite{2con}.  In this paper we have ruled out the possibility of black hole growth in the Randall-Sundrum model.  We therefore do not expect strange stars to collapse due to TeV gravity black hole seeding.  We have also presented formalisms for the growth of 5D black holes in dense media which might be of interest in other areas such as cosmology and hadronic collider physics."
    },
    "0206/astro-ph0206028_arXiv.txt": {
        "abstract": "Fast X-ray timing can be used to probe strong gravity fields around collapsed objects and constrain the equation of state of dense matter in neutron stars. These studies require extremely good photon statistics. In view of the huge collecting area of its mirrors, XEUS could make a unique contribution to this field. For this reason, we propose to include a fast X-ray timing capability in the focal plane of the XEUS mirrors. We briefly outline the scientific motivation for such a capability. We compute some sensitivity estimates, which indicate that XEUS could provide better than an order of magnitude sensitivity improvement over the Rossi X-ray Timing Explorer. Finally, we present a possible detector implementation, which could be an array of small size silicon drift detectors operated out of focus. ",
        "introduction": "The X-rays generated in the inner accretion flows around black holes (BHs) and neutron stars (NSs) carry information about regions of the strongly curved space-time in the vicinity of these objects. This is a regime in which there are important predictions of general relativity still to be tested. High resolution X-ray spectroscopy and fast timing studies can both be used to diagnose the orbital motion of the accreting matter in the immediate vicinity of the collapsed star, where the effects of strong gravity become important. With the discovery of millisecond aperiodic X-ray time variability (QPO) from accreting BHs and NSs, and brightness burst oscillations in NSs, the Rossi X-ray Timing Explorer (RXTE, Bradt et al. 1993) has clearly demonstrated that fast X-ray timing has the potential to measure accurately the motion of matter in strong gravity fields and to constrain masses and radii of NSs, and hence the equation of state of dense matter.  Although the prime objectives of XEUS are to perform spectroscopy of faint X-ray sources to trace the origin and evolution of hot matter back to the early ages of the Universe (see Hasinger et al. 2000; Bleeker et al. 2000), the large collecting area required to meet these objectives could also be used for X-ray timing studies. In this paper, we briefly outline how XEUS could broaden its science by having a fast X-ray timing capability in the focal plane (section 2). The potential of X-ray timing studies is described in more details in two companion papers by Michiel van der Klis and Cole Miller. In section 3, we give some sensitivity estimates for XEUS. Finally, in section 4, we present a possible implementation for the fast X-ray timing capability. ",
        "conclusions": "Probing strong gravity fields, constraining the equation of state of dense matter, and more generally studying the brightest X-ray sources of the sky with fast X-ray timing could be achieved with XEUS with the addition of a fast X-ray timing capability in the focal plane. This would greatly extend the capabilities of XEUS at low cost, with little impact on the primary objectives of the mission, which remain the spectroscopy of the most distant X-ray sources of the Universe. Current detector technology meet the science requirements for the timing of the brightest X-ray sources. The most attractive detector implementation would consist of an array of small size silicon drift detectors operated out of focus."
    },
    "0206/astro-ph0206502_arXiv.txt": {
        "abstract": "We reconstruct a 3+1 formalism of general relativistic electromagnetism, and derive the equations of motion of charged particles in the pulsar magnetosphere, taking account of the inclination between the rotation axis and the magnetic axis. Apart from the previous works where space charge is evaluated by assuming the flow velocity being the light speed, we analyze particle motion in the polar cap, finding that gravity changes significantly its dynamics and the condition for acceleration. ",
        "introduction": "The origin of radio emission from pulsars remains a great mystery. One of the likely scenarios is that particles are accelerated along open magnetic field lines and emit $\\gamma$-rays which subsequently convert into electron-positron pairs under a strong magnetic field. A combination of the primary beam and pair plasma provide the radio emission mechanism \\cite{Mel}. The process of particle acceleration has been intensively studied in flat spacetime \\cite{FAS,AS,Mes81,Shi91,Shi95,Shi97}. The field-aligned electric field is driven by deviation of the space charge from Goldreich-Julian charge density \\cite{GJ}, which is determined by the magnetic field geometry. Therefore, general relativistic (GR) effects on the field geometry are crucial for the formation of the field-aligned electric field and particle acceleration.  Muslimov \\& Tsygan (1992) initiated the GR analysis of electromagnetic fields around a pulsar; they solved Maxwell equations on the assumption that the particles move with the light speed. Shibata showed, however, that their assumption of constancy of particle speed was not always true \\cite{Shi91,Shi97}: the particle motion and space charge show oscillatory behavior even on field lines curving toward the rotation axis if the assumed current density is sub-critical. Mestel (1996) extended Shibata's equations of motion to include GR effects, though his analysis is not complete. Recently Harding and Muslimov (2001) extended Muslimov-Tsygan's analysis to include the pair-creation dynamics.  In this paper we derive the basic equations more rigorously and generally: we depart Muslimov-Tsygan (and Harding-Muslimov) model by solving equations of motion for the particles, and correct Mestel's equation and extend it to include the inclination angle between the magnetic axis and the rotation angle. Next we solve the electric field together with the particle motion for the region near the magnetic pole just above the surface. It turns out that the effect of gravity on particle dynamics cannot be ignored, and hence the condition for acceleration is also modified. As pointed by Arons (1997), local deficit of space charge takes place on most of the field lines. In particular, particle dynamics on the ``away'' field lines from the rotation axis is changed qualitatively by gravity if the current density is less than the critical value. It is also true that oscillatory solution exists even with GR effect. We find that some oscillatory solutions disappear above a certain height.  This paper is organized as follows. In section 2, we reconstruct a 3+1 formalism of GR electromagnetism. In section 3, we derive equations of motion of charged particles around a magnetized rotating star. In section 4, solving the equations of motion, we analyze particle acceleration in the polar cap. Conclusions and discussions are given in section 5. We use the unit $c=1$. ",
        "conclusions": ""
    },
    "0206/astro-ph0206444_arXiv.txt": {
        "abstract": "We present the photometric properties of 1279 blue variable stars within the LMC. Photometry is derived from the MACHO database. The lightcurves of the sample exhibit a variety of quasi-periodic and aperiodic outburst behavior. A characteristic feature of the photometric variation is that the objects are reddest when at maximum outburst. A subset of 102 objects were examined spectroscopically. Within this subset, 91\\% exhibited Balmer emission in at least one epoch, in some cases with spectacular spectral variability. The variability observed in the sample is consistent with the establishment and maintenance of the Be phenomenon. ",
        "introduction": "The MACHO photometric database constitutes an unprecedented resource for the study of stellar variability. With coverage of over 20 million stars within the LMC alone, the MACHO project is superior to competitors in duration (seven years) and spatial coverage. A primary goal of the MACHO project was to detect microlensing events in/towards the Magellanic Clouds and the bulge of the Galaxy. One of the initial unexpected contaminants in the detection of microlensing events were blue variables provisionally termed ``bumpers'' \\citep{coo95}. These objects were found to brighten $\\Delta V$ $\\sim$ 0.$^m$2-0.4 over the course of fifty to several hundred days in a manner much like that expected from microlensing events. However these objects did not exhibit the achromaticity expected of microlensing.  In the present work we extract from the catalogue of MACHO variables (over 210000 objects) those stars which are in the vicinity of the main-sequence and exhibit large amplitude variability (see Section \\ref{sect:sel}). Section \\ref{sect:phot} presents an overview of lightcurve morphology and colour variations. Section \\ref{sect:spec} presents spectroscopic observations of a subset of 102 objects from our catalogue. In Section \\ref{sect:nature} we discuss the physical nature of the blue variables. The goal of the current paper is to present preliminary data on these objects for future studies to expand upon. ",
        "conclusions": "We have presented a summary of the properties of 1279 blue variables within the Large Magellanic Cloud. Variations of up to $\\Delta V \\sim$0.$^m$4 are observed in the form of rapid ($\\sim$50d) or prolonged (100-500d) eruptive events as well as fading events and long term (years) trends. The $V$$-$$R$ colours show a characteristic ``redder when brighter'' behaviour.  A subset of 102 stars were examined spectroscopically, in a limited number of instances over multiple epochs.  Many objects show spectacular photometric and spectroscopic variability. The majority (91\\%) of the spectroscopic sample were found to exhibit Balmer emission lines in at least one epoch (i.e.\\ Be stars).  This strongly suggests that the mechanism giving rise to the variability seen with in the sample of LMC stars is related to the Be phenomenon. Firstly, the photometric variability is similar to that observed in Galactic Be stars. Secondly, emission strength is seen to coincide with redder $V$$-$$R$ colour and maximum $V$ luminosity as expected from known properties of Galactic Be stars.  The outbursts observed in the blue variable population can therefore be understood as mass eject events from the central star leading to the establishment and maintenance of the Be disk. Through further study of this population could yield important insight into the physical process of mass loss and the confinement of material in the vicinity of star which underlies the Be star phenomenon. The role of metallicity in the Be phenomenon will be examined in a future paper where we plan to investigate properties of the blue variable population of the Small Magellanic Cloud."
    },
    "0206/astro-ph0206358_arXiv.txt": {
        "abstract": "We report the discovery of nineteen new \\HI\\ shells in the Southern Galactic Plane Survey (SGPS).  These shells, which range in radius from 40~pc to 1~kpc, were found in the low resolution Parkes portion of the SGPS dataset, covering Galactic longitudes $l=253\\arcdeg$ to $l=358\\arcdeg$. Here we give the properties of individual shells, including positions, physical dimensions, energetics, masses, and possible associations.  We also examine the distribution of these shells in the Milky Way and find that several of the shells are located between the spiral arms of the Galaxy.  We offer possible explanations for this effect, in particular that the density gradient away from spiral arms, combined with the many generations of sequential star formation required to create large shells, could lead to a preferential placement of shells on the trailing edges of spiral arms. Spiral density wave theory is used in order to derive the magnitude of the density gradient behind spiral arms.  We find that the density gradient away from spiral arms is comparable to that out of the Galactic plane and therefore suggest that this may lead to exaggerated shell expansion away from spiral arms and into interarm regions. ",
        "introduction": "\\label{sec:intro} \\HI\\ shells, as a class of objects, largely determine the structure, dynamics and evolution of the interstellar medium (ISM).  These massive objects, which are usually detected as voids in the neutral hydrogen (\\HI) emission, range in size from tens to hundreds of parsecs, and in some cases even kiloparsecs \\citep[e.g.][]{rand96,deblok00}.  It is believed that \\HI\\ shells are formed through the combined effects of stellar winds and supernovae, which input $10^{51}-10^{53}$ ergs of energy into the ISM, ionizing the neutral medium and sweeping up a massive expanding shell. Because of the significant energies involved, \\HI\\ shells may be the large, deterministic structures needed to power the turbulent cascade of energy seen in spatial power spectrum analyses of the Galaxy \\citep{spangler01}. In addition, with lifetimes on the order of tens of millions of years large \\HI\\ shells long outlive the radiative lifetimes of their parent \\HII\\ regions and supernova remnants and hence can be used as fossils to study the effects of star formation in the Galaxy.  In the nearby Large and Small Magellanic Clouds hundreds of \\HI\\ shells dominate the structure of the \\HI\\ \\citep{staveley-smith97,kim98}.  In the Milky Way, however, the number of cataloged shells is considerably smaller. Existing shell catalogs, for example the \\citet{heiles79,heiles84} catalogs focus on the Northern Galaxy and are limited by undersampled, low resolution \\HI\\ surveys.  There are no comparable \\HI\\ shell catalogs for the Southern Galaxy.  Additionally, though there have been many studies of individual supershells \\citep[e.g.][]{maciejewski96,heiles98}, there are only a few studies of the global properties or distributions of Galactic \\HI\\ shells \\citep[e.g.][]{ehlerova96}.  Clearly Galactic \\HI\\ shell catalogs are in need of updating and expansion to more global scales.  Recent \\HI\\ surveys, such as the Southern and Canadian Galactic Plane Surveys \\citep[SGPS and CGPS;][]{mcgriff01a,taylor99}, are increasing the number of shells known in the Milky Way.  These surveys, which image the \\HI\\ in large regions in the Galaxy, benefit from coverage of spatial scales ranging from parsecs to kiloparsecs, enabling more complete catalogs of Galactic \\HI\\ shells.  With these surveys we hope to be able to determine the nature of many \\HI\\ shells and the extent to which the ISM in the Milky Way is shaped by them.  Here we present nineteen new \\HI\\ shells found in the Southern Galactic Plane Survey.  This catalog is limited to shells larger than one degree in angular diameter.  The shells range in physical size from 40 pc to 700 pc and are distributed throughout the Parkes SGPS region ($253\\arcdeg \\leq l \\leq 358\\arcdeg$; $|b| \\leq 10\\deg$).  We focus here on the properties of these shells and their distribution in the Galaxy.  In \\S \\ref{subsec:criteria} we discuss the selection criteria for the current catalog.  The observed and calculated properties of the shells are given in \\S \\ref{subsec:props} and several individual shells are described briefly in \\S \\ref{subsec:individ}. Selection biases in the catalog are discussed in \\S \\ref{subsec:biases}. In \\S \\ref{sec:distribution} we focus on the Galactic distribution of \\HI\\ shells, including those from the catalogs of \\citet{heiles79,heiles84}.  We explore shell properties in the context of spiral structure, investigating the density structure of the \\HI\\ disk (\\S \\ref{subsec:shells&spiral}).  Here we suggest that because of a density gradient from spiral arms to interarm regions, shells on the edges of spiral arms should attain larger sizes than shells in other regions. ",
        "conclusions": "\\label{sec:catconcl} We have discovered nineteen new \\HI\\ shells in the Southern Galactic Plane Survey.  These shells vary in radius from 40 pc to 700 pc, have expansion velocities between 6~\\kms\\ and 20~\\kms, have expansion energies between $10^{51}$ ergs and $10^{53}$ ergs and are distributed throughout the fourth quadrant of the Galaxy.  We have used this new catalog, along with those of \\citet{heiles79,heiles84}, to examine the distribution of large \\HI\\ shells in the Milky Way.  We find a tendency for large, energetic shells to be located at large Galactocentric radii; a tendency that has also been observed in external galaxies \\citep[e.g.][]{deul90,walter99}.  We also show that many shells are located between the spiral arms of the Galaxy.  We have used basic timescale arguments to show that the timescales are such that gas moves out of a spiral arm in a time that is comparable to the time required for the formation of a large ($R>100$ pc) shell.  Hence, we believe that large shells will often end their lives between spiral arms.  We have used the same timescale arguments to show that in the inner Galaxy, where few large shells are observed, a shell would move out of an arm, through the interarm region, and be struck by the next spiral arm before it could grow very large.  This may explain why large shells are only seen at large Galactic radii.  We have also used spiral density wave theory to explore the radial density structure of the Galactic disk as a result of spiral arms.  We compared this to the density structure extending out of the Galactic plane towards the halo to show that the density gradient away from spiral arms is comparable to that from the disk to the halo.  Because simulations have shown that shells expanding out of the plane tend to elongate in that direction, we suggest an analogous scenario in which a shell experiences runaway expansion away from spiral arms.  This effect, combined with the multiple generations of star formation required to create large shells, should lead to a population of interarm shells."
    },
    "0206/astro-ph0206114_arXiv.txt": {
        "abstract": "The origins of recently reported anisotropy of the local velocity field of nearby galaxies (velocities $<$ 500 km/s corresponding to the distance less than 8 Mpc) are studied. The exact solution of the Newtonian equation for the expanding Universe is obtained. This solution allows us to separate the Newtonian motion of nearby galaxies from the Hubble flow by the transition to the conformal coordinates. The relation between the Hubble flow and the Newtonian motion is established. We show that the anisotropic local velocity field of nearby galaxies can be formed by such a Newtonian motion in the expanding Universe, if at the moment of the capture of galaxies by the central gravitational field. ",
        "introduction": "%    Recent observation of the local velocity field of galaxies gives a three-dimensional ellipsoid with different values of the Hubble parameter, clearly showing its  anisotropic character ~\\cite{Karach1,Karach2}.  In this paper, we present a possible point of view that this local velocity field of galaxies can be explained by their Newtonian motions.  The analysis of the observational data will be based on the radial velocities of nearby galaxies, belonging to the Local Group. Our paper is organized in the following manner. In Section II, the cosmic evolution is described.  In Section III, we introduce the Newtonian motion and separate this motion from the cosmic one. In  Section IV, the initial data of the galaxies capture by a central gravitational field is considered.  In  Sections V, the simplest example is given to elucidate our results. The paper ends with the conclusions. ",
        "conclusions": "Our paper was motivated by the finding that in the local Universe the velocity field is anisotropic \\cite{Karach1,Karach2}. This effect is difficult for explanation. The only possible suggestion, but rejected by Karachentsev, was rotation \\cite{Karach2}. We are trying to find the origin of this anisotropy. In ordered to do this, we consider the general uniform expansion of the Universe. Since this is the nearby (less than 8 Mpc) part of the Universe around us, we use the Newtonian approach. We studied the motion of the test massive particle in the central gravitational field on the background of cosmic evolution of the type of FLRW space--time with uniform expansion. We assume the rigid state of the matter when densities of energy and pressure are equal and which corresponds to conformal cosmology \\cite{039} compatibles with Supernova data \\cite{snov}. We obtained the exact solution of the above--mentioned Kepler problem, to find the difference between the uniform Hubble flow and our case. We have shown that this difference was anisotropic. In such a way we explained the anisotropy of the local velocity field by the Newtonian motion of galaxies in the central field. Of course, our 2--D consideration  shows a possible mechanism of the observed 3--D anisotropy.  Having the solution of the Kepler problem we admit the rotation of galaxies around the center of the local Universe. This local Universe must be regarded as the Local Group of galaxies. In such a way, we support the picture in which galaxies rotate around the center of the Local Group in the class of ellipsoidal trajectories."
    },
    "0206/astro-ph0206322_arXiv.txt": {
        "abstract": "We study the effects of the outer boundary conditions in neutrino-driven winds on the r-process nucleosynthesis.  We perform numerical simulations of hydrodynamics of neutrino-driven winds and nuclear reaction network calculations of the r-process.  As an outer boundary condition of hydrodynamic calculations, we set a pressure upon the outermost layer of the wind, which is approaching toward the shock wall.  Varying the boundary pressure, we obtain various asymptotic thermal temperature of expanding material in the neutrino-driven winds for resulting nucleosynthesis.  We find that the asymptotic temperature slightly lower than those used in the previous studies of the neutrino-driven winds can lead to a successful r-process abundance pattern, which is in a reasonable agreement with the solar system r-process abundance pattern even for the typical proto-neutron star mass $M_{NS} \\sim 1.4 M_{\\odot}$.  A slightly lower asymptotic temperature reduces the charged particle reaction rates and the resulting amount of seed elements and lead to a high neutron-to-seed ratio for successful r-process.  This is a new idea which is different from the previous models of neutrino-driven winds from very massive ($M_{NS} \\sim 2.0$ $M_{\\odot}$) and compact ($R_{NS} \\sim 10$ km) neutron star to get a short expansion time and a high entropy for a successful r-process abundance pattern.  Although such a large mass is sometimes criticized from observational facts on a neutron star mass, we dissolve this criticism by reconsidering the boundary condition of the wind.  We also explore the relation between the boundary condition and neutron star mass, which is related to the progenitor mass, for successful r-process. ",
        "introduction": "The r-process nucleosynthesis, which is a rapid neutron capture process faster than beta-decay, is believed to be responsible for about a half of elements heavier than iron (Burbidge et al. 1957).  Many heavy elements presumably produced in the r-process have recently been detected in extremely metal-poor stars by recent astronomical observations (Sneden et al. 1996, 1998, 2000, Westin et al. 2000, Johnson and Bolte 2001, Cayrel et al. 2001, Honda 2001).  Their abundance pattern proves to be very similar to the one of the solar r-process pattern, which is called the universality in the r-process abundance pattern, especially in the region of $56 \\le Z \\le 70$.  This universality of abundance pattern strongly suggests that the r-process occurs in the same way independently of the metallicity along the entire history of Galactic chemical evolution from the beginning of the Galaxy to the present. It means simultaneously the fact that the origin of the r-process is most likely in  supernovae (SNe) of massive progenitor stars since massive stars first evolve and end up with SN explosions whose ejecta reflects  abundances in metal-poor stars.  One of the most plausible sites of the r-process is the neutrino-driven winds in supernovae.  Woosley et al. (1994) have demonstrated that the very high entropy conditions, $\\sim 400$ k$_B$, are realized in the neutrino-driven wind and on these specific conditions the r-process nucleosynthesis occurs successfully.  This r-process scenario in high entropy hot bubble has been later pointed out to be rather difficult because copious supernova neutrinos change neutrons to protons and hinder the r-process (Meyer 1995).  In more recent studies, the successful r-process pattern in the neutrino-driven winds has been obtained even for relatively low entropy, $\\sim 200$k$_B$, provided that the expansion time scale is much shorter, $\\sim 10$ ms, than the time scale of the neutrino process (Otsuki et al. 2000, Sumiyoshi et al. 2000), making the neutrino-driven wind scenario viable again.  However, they assumed a large neutron star mass, $\\sim 2.0 M_{\\odot}$, in order to gain a slightly higher entropy.  These parameter setups have been referred with caution because the observed neutron star has a typical mass $\\sim 1.4 M_{\\odot}$ and a radius $\\sim 10$ km.  In this paper we will discuss that the outer boundary conditions of the neutrino-driven winds may resolve this problem of the neutron star mass. It has recently been founded (Terasawa et al. 2001) that the r-process occurs far from the neutron star in a less dense region rather close to and behind the outward shockwave after the surface material is blown off.  It is generally known that alpha captures are very sensitive to the temperature. A little change in temperature can make a large effect on the abundance of seed elements and heavy r-process elements. This effect on nucleosynthesis of outer region has not been studied very well.  In most previous studies, the boundary condition of outer region has been chosen based on the results by Woosley et al. (1994). This is because their simulation is the unique simulation which deals with both supernova explosion and neutrino-driven wind at once. In their study the temperature is about 0.1 MeV and the density is $\\sim 10^{3}$ g/cm$^3$ at the outer boundary (at radius of 10$^{4}$ km).  The physical condition in the outer boundary region depends on the competition between the falling matter from the envelope and the outward shockwave, giving many factors to change by a core bounce. There are a lot of factors to effect the boundary condition.  We perform the hydrodynamical simulations of neutrino-driven winds and investigate the dependence of the r-process nucleosynthesis on the pressure of the outer material, toward which the wind blows, behind the shockwave.  We show that the successful r-process can occur in the neutrino-driven winds from a typical neutron star mass with $1.4 M_{\\odot}$ by choosing a slightly lower pressure and temperature than those adopted previously.  We also examine the dependence of the r-process yields on the neutron star mass in order to assess the relation between the stellar mass and the outer boundary condition. Since a different mass of progenitor star is suggested to lead to a different mass of envelope and remnant mass, we infer a correspondence between the progenitor mass and the outer boundary condition to realize the universality and successful the r-process abundance pattern.  This study is important to constrain the mass range of progenitor stars that culminate their evolution for collapse-driven supernovae as the source of the r-process elements which are observed in very metal-poor stars. ",
        "conclusions": "We showed a possibility that the r-process can successfully occur in the neutrino-driven winds from a neutron star having typical observed neutron star mass, $1.4 M_{\\odot}$, provided that the outer boundary condition is appropriately chosen.  More specifically from Fig. 1, we can conclude that the model with $P_{out} = 10^{20}$ dyn/cm$^2$ (dashed line) is most likely in the neutron star model with $M_{NS} =1.4 M_{\\odot}$ and $R_{NS} = 10$ km.  In the present model the dynamical timescale is a few times longer than the previous successful models and the entropy per baryon is relatively low, $100-200$ k$_{\\rm B}$.   This conclusion results from the fact that the alpha captures do not efficiently work as the outer boundary temperature, $T_{out}$, becomes lower.  We can reconfirm this effect quantitatively in the following way.  If the seed nuclei are synthesized only by the alpha-process and the seeds do not change to other heavy nuclei approximately, the decrease in alpha particles from the peak value in Fig. 1 is equal to the increase in the seeds. Then, the relation, $\\Delta Y_{\\alpha} \\times 4 = \\Delta Y_{seed} \\times 100$, follows due to the mass conservation, where we assume the averaged mass number of seed nuclei 100.  We could confirm actually that this relation holds very well in final abundances ($Y_{seed,out}$ and $Y_{\\alpha,out}$) in our three models (Table 1).   It is generally believed that the neutron star mass depends somewhat on the progenitor mass (Woosley \\& Weaver 1995, Thielemann et al. 1996, Timmes et al. 1996, Limongi et al. 2000).  The physical conditions that govern the r-process are determined by $S$, $Y_e$, and $\\tau_{dyn}$, which also strongly depend on the neutron star mass. However, recent observations of neutron-capture elements in metal-deficient stars show the universality in the abundance pattern for the r-process elements with $56 \\le Z \\le 70$ (Sneden et al. 2000, Westin et al. 2000, Cayrel et al. 2001, Honda 2002). They suggest that the universal r-process abundance pattern should be realized independent of the neuron star mass.  Let us consider the dependence of abundance pattern on the neutron star mass by using the same outer boundary pressures.  We adopt several neutron star mass of $1.2$, $1.3$, $1.4$, $1.5$ and $1.6 M_{\\odot}$ in the simulations of the neutrino-driven winds (Terasawa 2002 and Terasawa et al. 2002). These masses cover the almost all observed neutron star masses except for a few neutron stars (Bulik et al. 1995, Brown et al. 1996, Thorsett and Chakrabarty 1999, Barziv et al. 2001, and references therein).  As for the outer boundary conditions, we varied the pressure values as $P_{out} = 10^{20}$, $10^{21}$, and $10^{22}$ dyn/cm$^2$ in our simulations of the $1.2 - 1.6 M_{\\odot}$ models. The value of $S$ is higher and $\\tau_{dyn}$ is shorter as the neutron star mass becomes larger, although $Y_{e,i}$ are almost the same $\\sim 0.43 - 0.44$ in all calculations, which we have carried out, and $T_{out}$ is common for each value of $P_{out}$. As discussed in the previous section, we saw that $Y_{\\alpha,out}$ becomes smaller and $Y_{seed,out}$ becomes larger with increasing pressure for the fixed neutron star mass. On the other hand, for the fixed outer boundary pressure, $Y_{\\alpha,out}$ becomes larger and $Y_{seed,out}$ becomes smaller as the neutron star mass becomes larger. From these systematics which we found, we understand that more abundant heavy elements are synthesized as the neutron star mass is larger and the outer boundary pressure becomes lower.  In the previous study of the r-process nucleosynthesis in neutrino-driven winds (Sumiyoshi et al. 2001), the outer boundary pressure was set equal to $P_{out} \\sim 10^{22}$ dyn/cm$^2$ for both neutron star mass models of $1.4 M_{\\odot}$ and $ 2.0 M_{\\odot}$. This pressure ($P_{out} \\sim 10^{22}$ dyn/cm$^2$) corresponds to the temperature $T_{out} \\sim 0.1$ MeV.  By their calculations, when this value was adopted in the $1.4 M_{\\odot}$ model, the flow of nucleosynthesis virtually stopped at the nuclear mass region below the 2-nd peak.  Therefore, a higher neutron star mass model ($2.0 M_{\\odot}$) was adopted to increase the entropy for a successful r-process. This theoretical correlation among $P_{out}$, $T_{out}$, and $M_{NS}$ is reasonably understood, that is a higher outer boundary pressure corresponds to a higher neutron star mass, and vice versa. However, it is to be stressed that increasing boundary pressure and temperature leads to unsuccessful r-process without increasing the neutron star mass.  In the present study we find another condition to realize successful r-process namely by adopting suitable outer boundary condition with keeping the neutron star properties as those measured observationally, i.e. $M_{NS} \\sim 1.4 M_{\\odot}$ and $R_{NS} \\sim 10$ km.  This result makes a strong constraints on modeling the dynamics of supernova explosion in view of constructing the structure model of massive progenitor stars in order to clarify the physical conditions of the r-process nucleosynthesis.    \\begin{deluxetable}{ccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{The Key Quantities for the \\lowercase{r}-Process Nucleosynthesis.  \\label{tbl-1}} \\tablewidth{0pt} \\tablehead{ \\colhead{P$_{out}$ [dyn/cm$^2$]} & \\colhead{$\\tau_{dyn}$ [sec]}   & \\colhead{ S [k$_B$]}   & \\colhead{$Y_{e,i}$} & \\colhead{T$_{out}$ [$10^9$K]}  & \\colhead{$Y_{\\alpha,out}$} & \\colhead{$Y_{seed,out}$}} \\startdata  $10^{20}$   &  $2.32 \\times 10^{-2}$   & $200$ & 0.43 & 0.4 &$0.196$&$9.2 \\times 10^{-4}$\\\\ $10^{21}$   &  $2.54 \\times 10^{-2}$   & $180$ & 0.43 & 0.7 &$0.189$&$2.0 \\times 10^{-3}$\\\\ $10^{22}$   &  $3.34 \\times 10^{-2}$   & $170$ & 0.44 & 1.3 &$0.174$&$2.4 \\times 10^{-3}$\\\\  \\enddata   \\tablecomments{The summary of model parameters and key quantities for the r-process. P$_{out}$ is the pressure at the outer boundary and we give these values as an outer boundary condition. $\\tau_{dyn}$, $S$, and $Y_{e,i}$ are the dynamical timescale, entropy per baryon, and initial electron fraction, respectively. Note that the definition of $\\tau_{dyn}$ is the $e$-fold time at $T = 0.5$ MeV. $Y_{\\alpha,out}$ and $Y_{seed,out}$ are the final abundances of alpha-particles and seed nuclei.} \\end{deluxetable}   \\begin{figure*} \\includegraphics[width=\\textwidth]{f1.eps} \\caption{Time variations of the alpha-particle abundance, $Y_{\\alpha}$, and temperature $T_9$ (upper panel), and neutron-to-seed ratio, $Y_n/Y_{seed}$, and seed abundance,$Y_{seed}$ (lower panel). The dashed, solid, dotted lines are the results in the cases of $P_{out} = 10^{20}$, $10^{21}$, and $10^{22}$ dyn/cm$^2$, respectively. \\label{ntos}} \\end{figure*}    \\begin{figure*} \\includegraphics[width=\\textwidth]{f2.eps} \\caption{Final abundance yields as a function of the mass number. The dashed, solid, dotted lines are the same as those in Fig. 1. Data points are the solar r-process abundances in arbitrary unit from K$\\ddot{\\rm a}$ppeler et al. (1989). \\label{final}} \\end{figure*}"
    },
    "0206/astro-ph0206052_arXiv.txt": {
        "abstract": "}[1]{{ \\footnotesize \\noindent {\\bf Abstract:} #1 \\\\}} \\renewcommand{\\author}[1]{\\subsubsection*{\\it#1}} \\newcommand{\\address}[1]{\\subsubsection*{\\it#1}} \\newcommand{\\ltsimaZaroubi}{$\\; \\buildrel < \\over \\sim \\;$} \\newcommand{\\lsimZaroubi}{\\lower.5ex\\hbox{\\ltsimaZaroubi}} \\newcommand{\\gtsimaZaroubi}{$\\; \\buildrel > \\over \\sim \\;$} \\newcommand{\\gsimZaroubi}{\\lower.5ex\\hbox{\\gtsimaZaroubi}}  \\chapter*{Cosmic Flows: Review of Recent Developments}  \\author{Saleem Zaroubi} \\address{Max Planck Institute for Astrophysics\\\\ Karl-Schwarzschild-Str. 1\\\\ D-85741 Garching, Germany }  \\abstract{I review the recent developments in the analysis of cosmic flow data, in particular, latest results of bulk flow measurements, comparison between redshift and peculiar velocity catalogs with emphasis on the measured value of the $\\beta\\, (=\\Omega_m^{0.6}/b)$ parameter, and matter power spectrum estimates from galaxy peculiar velocity catalogs. Based on these developments, one can argue that most of the previous discrepancies in the interpretation of cosmic flow data, {\\it maybe} with the exception of bulk flow measurements on scales $\\gsimZaroubi 100 h{^{-1}}{\\rm Mpc}$, have either been resolved or fairly understood.  } ",
        "introduction": "\\label{sec:zaroubi_introduction} Within the gravitational instability (GI) framework for the growth of cosmic structures, the peculiar velocity field of galaxies and clusters provides a direct and reliable probe of the matter distribution, under the natural assumption that these objects are unbiased tracers of the large-scale, gravitationally induced, velocity field. The GI paradigm requires that the linear peculiar velocity, ${\\bf v}({\\bf r})$ -- defined as the deviation from Hubble expansion -- and linear mass density contrast, $\\delta_m({\\bf r})$, be related to one another according to the local (differential) relation, \\begin{equation} \\nabla \\cdot {\\bf v} = -\\Omega_m^{0.6} \\delta_m, \\label{eq:zaroubi_divv} \\end{equation} or its global (integral) counterpart, \\begin{equation} {\\bf v} = {\\Omega_m^{0.6} \\over 4\\pi} \\int d^3{\\bf r'} { \\delta_m({\\bf r'}) ({\\bf r'}-{\\bf r}) \\over \\vert {\\bf r'}-{\\bf r}\\vert^3}, \\label{eq:zaroubi_intd} \\end{equation} where $\\Omega_m$ is the matter overdensity parameter. Note that the peculiar velocity field is determined by the distribution of the matter with all its components especially the dominant dark matter component.  In order to measure peculiar velocities of galaxies and clusters, observers use a variety of distance indicators. Generally, these indicators relate two quantities, one among which is distance dependent, {\\it e.g.,} galaxy luminosity, and the other is distance independent, {\\it e.g.,} galaxy rotational velocity. The best known examples of such indicators are the Tully-Fisher~\\cite{tully77} and Faber-Jackson~\\cite{faber76} relations; over the last decade or so these and many other types of distance indicators have been used to measure cosmological distances. The availability of a large number of galaxy peculiar velocity catalogs, some of them with few thousands objects, have turned cosmic flows to one of the main probes used to study the large scale structure in the nearby universe. Here I'll concentrate on the following three statistical measures of the velocity field:  \\begin{enumerate} \\item {\\it The Bulk Flow:} This measure, defined as the average streaming motion within certain volume, is probably the easiest statistic to estimate from the observed radial component of peculiar velocities. At the Cosmic Microwave Background radiation (CMB) restframe, the bulk motions are expected to converge to zero with increasing volume. The rate of convergence depends on the fluctuations in the matter distribution on various scales, {\\it i.e.,} the large scale matter fluctuations power spectrum. This dependence on cosmological models has motivated several attempts to measure the dipole component of the local peculiar velocity field and to determine the volume within which the streaming motion vanishes. As of yet, bulk flow measurements have produced conclusive and consistent results only on scales $\\lsimZaroubi 60h{^{-1}}{\\rm Mpc}$, but failed to do so on scales $\\gsimZaroubi 100 h{^{-1}}{\\rm Mpc}$ (for recent works see \\cite{colless01,dacosta00, courteau00, dale99, giovanelli97, hudson99, willick99}).  \\item {\\it The mass power spectrum:} Equation~\\ref{eq:zaroubi_intd} suggests that one can estimate the bias free, $\\Omega_m^{1.2}$ weighted, matter power spectrum directly from the measured peculiar velocities. To date, likelihood analysis based estimations of the matter power spectrum exist for the Mark~III~\\cite{zaroubi97}, SFI~\\cite{freudling99} and ENEAR~\\cite{zaroubi01} catalogs. All of these measurements has consistently produced power spectra with amplitudes larger than those measured by other data sets, {\\it e.g.,} galaxy redshift surveys.  \\item {\\it The $\\beta$ parameter:} Peculiar velocities enable a reconstruction of the large scale matter distribution independent of redshift surveys. Therefore, one can use Eqs.~\\ref{eq:zaroubi_divv} and~\\ref{eq:zaroubi_intd} to compare the matter and 3D velocity distributions deduced from the measured radial peculiar velocities to those obtained from redshift surveys. This comparison requires biasing model which specifies how galaxies follow the underlying {\\it total} matter distribution. On the scales of interest, it is usually assumed that \\begin{equation} \\delta_g = b \\delta_m , \\label{eq:zaroubi_bias} \\end{equation} where $\\delta_g$ is the galaxy observed density contrast and $b$ is the linear bias parameter.  The comparison is used to: 1) test the validity of Eqs.~\\ref{eq:zaroubi_divv} and~\\ref{eq:zaroubi_intd}, namely, the basic GI paradigm; 2) test the linear biasing model; 3) directly measure the value of $\\Omega_m$ (\\cite{branchini01, dacosta98, davis96, sigad98, willick98, willick97b} and \\cite{zaroubi02b}). Until recently, comparisons using Eq.~\\ref{eq:zaroubi_divv} have systematically yielded $\\beta$ values larger than those obtained from using Eq.~\\ref{eq:zaroubi_intd}. \\end{enumerate}  As mentioned above, many have attempted to estimate these statistical measures of peculiar velocities over the years. Until recently, the results they obtained have often been inconsistent with each other and with estimations from other data, depending on the specific peculiar velocity catalog at hand and the analysis methods.  It is generally accepted that the main reason for the inconsistencies lies in the problematic nature of the distance indicators used to determine the peculiar velocities. First, the distance measurements carry large random errors, including intrinsic scatter in the distance indicator and measurement errors, which grow in proportion to the distance from the observer and thus become severe at large distances. Further nontrivial errors are introduced by the nonuniform sampling of the galaxies that serve as velocity tracers. In particular, the Galactic disk obscures an appreciable fraction of the sky, creating a significant``zone of avoidance'' of at least $40\\%$ of the sky. When translated to an underlying smoothed field, these errors give rise to severe systematic biases. In light of these difficulties, the inconclusive results have led many to question the reliability of the peculiar velocity datasets as cosmological probes.  \\begin{figure} \\vskip -1.5 truecm \\centering \\begin{tabular}{c} \\hskip -1 truecm \\mbox{\\epsfig{file=zaroubi_fig1a.eps,height=9cm}} \\\\ \\mbox{\\epsfig{file=zaroubi_fig1b.eps,height=7cm}} \\end{tabular} \\vskip 1 truecm \\caption{{\\small Bulk Flow measurements. Upper panel: the symbols show the amplitude of the measured bulk flow (with its error) from the following surveys: Surface Brightness Fluctuations (SBF), SFI , ENEAR (EN), Shellflow (SF), Supernovae type Ia (SNIa), SMAC, EFAR, LP10, SCII and LP (see table for explanation) as a function of radius. The CMB dipole COBE measurement and bulk flow from the PSC$z$ redshift catalog are also shown. The solid line shows the expected rms bulk velocity of a sphere of radius $R$ for standard $\\Lambda$CDM model; the dashed lines represent 1-$\\sigma$ cosmic scatter about the rms.  Lower panel: the symbols show the direction of some of the measured bulk flow vectors, note that the catalogs that correspond to $R\\sim 60 h{^{-1}}{\\rm Mpc}$ have consistent directions while measurements that correspond to large distances do not. }} \\label{fig:zaroubi_fig1} \\end{figure} In this article, I review the recent developments in these three areas and show that significant improvements have occurred in the field in the past few years. In addition, an argument is put forward that those developements lead to alleviating most of the inconsistencies and to understanding, at least qualitatively, the cause of the remaining outstanding ones. ",
        "conclusions": "\\label{sec:zaroubi_summary}  To evaluate the outstanding issues in the peculiar velocity measurements in general terms one should differentiate between two types of problems. The first is the inconsistencies among results obtained from various peculiar velocity data sets and occasionally even within those obtained from the same data set. The second, is the disagreement with measurements from other types of data sets. The former is obviously more severe as it implies unresolved systematics in the peculiar velocity datasets themselves.  Measuring by this yard stick, the conflicting results obtained from the bulk flow and the $\\beta$ parameter measurements are more serious than the higher amplitude found with the power spectrum measurements.   It is reassuring that within the local universe ($R \\lsimZaroubi 60 h{^{-1}}{\\rm Mpc}$) all the recent bulk flow measurements, especially from the SFI, Shellflow and ENEAR catalogs do agree very well with each other both in terms of amplitude and direction. Those three marginally agree with the Mark~III data, which gives a somewhat higher value of $V_b$. Given that the Shellflow sample have clearly shown that the Mark~III data set is somewhat miscalibrated, this disagreement is hardly an issue.  Unfortunately however, the picture on larger scales ($R\\gsimZaroubi 100 h{^{-1}}{\\rm Mpc}$) is very different and the disagreement among the various measurements is yet to be resolved. While the bulk flow obtained from the SNIa~\\cite{riess97}, SCI/SCII~\\cite{dale99} and EFAR~\\cite{colless01} samples clearly points toward convergence of the CMB dipole; the SMAC\\cite{hudson99}, LP \\cite{lauer94} and LP10 \\cite{willick99} surveys find a bulk flow amplitude of $\\sim 600-700 {\\rm km\\,s{^{-1}}}$. Although the former three measurements are consistent with the one obtained from the PSC$z$ redshift catalog data, the NVSS radio sources and with theoretical prejudice, by no means the later three have been refuted. It is worth pointing out however, that the LP bulk flow, while comparable in amplitude, disagrees with the SMAC and LP10 flow. Furthermore, almost all of the deep peculiar velocity surveys have small number of objects and therefore probably prone to the pitfalls of small number statistics.   The peculiar velocity measurements have systematically led to mass power spectrum amplitudes higher than those obtained from other types of data. In light of the overwhelming evidence pointing towards lower amplitude power spectrum, a special effort should be made to rule out any inherent bias in the {\\it prior} assumptions made in the peculiar velocity based power spectrum measurements. In fact, two recent studies~\\cite{hoffman00, silberman01} strongly suggest a generic problem with the theoretical framework used to estimate the mass power spectrum from the Mark~III, SFI, and ENEAR velocity surveys.  Possible sources of this problem lie with non-linear dynamical effects and/or oversimplified treatment of errors~\\cite{silberman01}.    The irreconcilability of the $\\beta$ estimation from density-density and velocity-velocity comparisons has been one of the major outstanding issues in the cosmic flows field of study. A new technique, the UMV, have enabled carrying out, for the first time, both comparisons in the same framework. The results obtained by applying this technique to several galaxy peculiar velocity catalogs yield low values of the $\\beta$ parameter ($\\sim 0.5-0.6$), a result consistent with those obtained from the previous velocity-velocity comparisons and from the analysis of redshift surveys.  These latest results clearly strengthen the case for low $\\beta$ values, in agreement with those obtained by the previous velocity-velocity studies. A qualitative inspection into the reason of the discrepancy in the estimation of $\\beta$ between the UMV and POTENT density-density comparisons allows one to speculate that the most likely explanation is a collusion between both the systematic errors in the Mark~III data and the noise effects on the POTENT algorithm, that somehow conspired to produce these high $\\beta$-values.  In light of the developments presented in this review, it is argued that most of the outstanding issues in the large scale peculiar velocity field of study, {\\it maybe} with the exception of the large scale velocity field, have been either resolved or understood (at least qualitatively) and we finally have a consistent cosmological model emerging from the study of cosmic flows.  The experience gained during the convergence towards the current status, despite the crooked path it took, will be invaluable when dealing with the large future datasets.   \\bigskip  \\noindent {\\Large{\\bf Acknowledgements}}  \\bigskip  Much of the work presented here have been done in collaboration with E. Branchini, L.N. da Costa and Y. Hoffman, their contribution is acknowledged.  I would like to thank the organizers of this meeting for suggesting to me to review this topic."
    },
    "0206/astro-ph0206264_arXiv.txt": {
        "abstract": "We report the first detection of an \\ion{O}{8} Ly$\\alpha$ absorption line associated with an overdense region in the intergalactic medium (IGM) along the sightline towards PKS~2155-304 with the {\\sl Chandra} Low Energy Transmission Grating Spectrometer (LETGS). The absorption line is detected at $4.5\\sigma$ level with $cz \\approx 16,600~\\rm km\\ s^{-1}$. At the same velocity \\citet{sps98} detected a small group of spiral galaxies (with an overdensity of $\\delta_{gal} \\sim 100$) and low metallicity \\ion{H}{1} Ly$\\alpha$ clouds.  We constrain the intragroup gas that gives rise to the \\ion{O}{8} Ly$\\alpha$  line to a baryon density in the range $1.0 \\times 10^{-5} < n_{b} < 7.5\\times 10^{-5}~\\rm cm^{-3}$ ($50 < \\delta_{b} < 350$) and a temperature of $4-5\\times10^{6}$ K, assuming 0.1 solar abundance. These estimates are in accordance with those of the warm/hot intergalactic medium (WHIM) that are predicted from hydrodynamic simulations.  Extrapolating from this single detection implies a large fraction of the ``missing baryons'' ($\\sim$ 10\\%, or $\\sim$ 30-40\\% of the WHIM) are probed by the \\ion{O}{8} absorber. ",
        "introduction": "The cosmic baryon budget at low and high redshift indicates that a large fraction of baryons in the local universe have so far escaped detection (e.g., \\citealp{fhp98}). While there is clear evidence that a significant fraction of these ``missing baryons'' (between 20-40\\% of total baryons) lie in photoionized, low-redshift Ly$\\alpha$ clouds \\citep{pss00}, the remainder could be located in intergalactic space with temperatures of $10^{5}-10^{7}$ K (warm-hot intergalactic medium, or WHIM). Resonant absorption from highly-ionized ions located in the WHIM gas has been predicted based on both analytic studies of structure formation and evolution (\\citealp{sba81,aem94,plo98,fca00}) and cosmic hydrodynamic simulations (\\citealp{hgm98,cos99a,dco01,fbc02}). Recent discovery of \\ion{O}{6} absorption lines by the Hubble Space Telescope ({\\sl HST}) and the Far Ultraviolet Spectroscopic Explorer ({\\sl FUSE}) (\\citealp{tsa00, tsj00, tgs01}) indicates that there may be a significant reservoir of baryons in \\ion{O}{6} absorbers. While Li-like \\ion{O}{6} probes about $\\sim 30-40\\%$ of the WHIM gas (\\citealp{cto01, fbr01}), the remaining $\\sim 60-70\\%$ is hotter and can only be probed by ions with higher ionization potentials, such as H- and He-like Oxygen.  The WHIM gas, predicted by hydrodynamic simulations, typically has an overdensity of $\\delta_{b} = 5-200$ and is distributed in small groups of galaxies or in large scale filamentary structures that connect virialized halos. Given the expected physical conditions, current instruments such as {\\sl Chandra} and {\\sl XMM} should be capable of detecting resonance features from the WHIM gas \\citep{fbc02}. A first attempt by \\citet{fmb01} with the {\\sl Chandra} High Energy Transmission Grating Spectrometer (HETGS) yielded an upper limit of ${\\rm N(O~VIII) < 10^{17}\\ cm^{-2}}$. Recently, with the {\\sl Chandra} LETGS-HRC, \\citet{nic02} discovered X-ray resonant absorption features from warm/hot local gas along the line of sight towards PKS~2155-304, giving the first X-ray evidence for possible WHIM gas.  These absorbers are near zero redshift and therefore most likely  associated with the Galaxy or local group (see \\S~2).  We report the first detection of an \\ion{O}{8} Ly$\\alpha$ absorption line along the sightline towards PKS~2155-304 at $cz \\sim 16,600~\\rm km~s^{-1}$ plausibly associated with an {\\it intervening}, modestly overdense region in the IGM. The observations were made with the {\\sl Chandra} LETG-ACIS (Advanced CCD Imaging Spectrometer).  Physical diagnostics show the gas that gives rise to this resonance line has the typical properties of the WHIM, indicating the detection of a significant fraction of the ``missing baryons''. ",
        "conclusions": "Our detection of \\ion{O}{8} Ly$\\alpha$ from an intervening overdense region in the IGM at $cz \\sim 16,600\\rm~km~s^{-1}$, together with the detection by \\citet{nic02} of absorption from systems at $cz \\sim 0$ associated with the Galaxy or local group, have begun to reveal the much-anticipated warm-hot component of intergalactic matter (\\citealp{hgm98,cos99a,dco01,fbc02}).  The \\ion{O}{8} Ly$\\alpha$  feature we observe is plausibly associated with the same cosmic overdensity that gives rise to the \\ion{H}{1} Ly$\\alpha$ clouds and the four spiral galaxies found by \\citet{sps98}. Three of the galaxies have radial velocities within $\\sim100~\\rm km~s^{-1}$ of one another, while the fourth is within $\\sim400~\\rm km~s^{-1}$.  In contrast, the velocities of the \\ion{H}1 Ly$\\alpha$ clouds span $\\sim1800~\\rm km~s^{-1}$, with two falling very close to the velocities of the closest projected galaxies. The galaxies represent an overdensity of $\\delta_{gal} \\sim 100$ relative to the mean. Our estimates for the overdensity of the hot absorbing gas are $\\delta_{b} \\sim$50-350. The \\ion{H}{1} Ly$\\alpha$ clouds have very low metallicity, indicating that they are primarily composed of primordial material, whereas the hot gas should have an oxygen abundance of $\\sim 0.1$ solar to give sensible values for the column density. It is unlikely that the same gas is responsible for both \\ion{H}{1} and \\ion{O}{8} absorption. Following \\citet{sps98}, we identify this as a region of modest overdensity (the galaxy group) connected to filamentary or sheet-like structures demarcated by the \\ion{H}{1} Ly$\\alpha$ clouds, such as those seen in numerous cosmological simulations.  We note that the detection of a system such as this is roughly consistent with previous simulations.  Based on the models of \\citet{plo98} and \\citet{fbc02}, we estimate the cumulative absorption line number along a random sightline (Fig.~2). We adopt the numerical relation between baryon overdensity and metallicity from \\citet{cos99b} and assume collisional ionization. This observation is $\\sim 1.5\\sigma$ above the predicted value at the measured column density. According to Fig.~3, the probability of detecting such an \\ion{O}{8} Ly$\\alpha$ line or stronger to $z\\sim 0.1$ is $\\sim$ 10\\%. The consistency, if borne out by further studies, suggests that the simulations are valid descriptors of the warm/hot component of the IGM, at least for overdensities of $\\sim$100. Using those simulations, we estimate that we can probe about 10\\% of total baryons and about 30-40\\% of the WHIM gas in the local universe.  Other observations of bright AGN with the {\\sl Chandra} and {\\sl XMM-Newton} grating spectrometers are likely to reveal additional WHIM absorption features, and may turn the present few samplings into a real forest of high ionization Ly$\\alpha$ or He$\\alpha$ lines. However, the sensitivity of these instruments permit us only to probe the high density tail of the distribution (e.g. \\citealp{fca00}). It will be left to future missions like {\\sl Constellation-X} and  {\\sl XEUS}, assuming they have both high throughput and spectral resolving powers of 1000 or more at energies 0.1-1 keV, to fully reveal the X-ray forest.  We are grateful to F.~Nicastro and his colleagues for sharing an advance copy of their paper.  We thank the other members of the MIT/CXC team for its support. We also thank the referee, J.~M.~Shull, for helpful comments. This work is supported in part by contracts NAS 8-38249 and SAO SV1-61010.  \\clearpage"
    },
    "0206/astro-ph0206270_arXiv.txt": {
        "abstract": "{I present the various capabilities of upgraded and next generation radio telescopes, in particular their ability to detect and image distant star forming galaxies. I demonstrate that e-MERLIN, EVLA and LOFAR can detect systems similar to Arp 220 out to cosmological distances. The SKA can detect such systems out to any reasonable redshift that they might be expected to exist. Employing very long integration times on the multiple-beam SKA will require the array to be extended beyond the current specification - simply to avoid confusion noise limitations at 1.4~GHz. Other arguements for extending the SKA baseline length are also presented. As well as going ``deeper'' all these instruments (especially LOFAR and the SKA) will also go ``wider'' - detecting many tens of thousands of galaxies in a single day's observing. I briefly comment on the prospects of detecting radio emission at much earlier epochs, just before the epoch of re-ionisation.} ",
        "introduction": "The continuum sensitivity of radio astronomy instruments is set to improve by at least an order of magnitude over the next 10 years. Developments include the broad-banding of existing facilities (e.g. e-MERLIN and the EVLA) or the design and construction of entirely new, next generation instruments, such as the Low Frequency Array (LOFAR) and the Square Km Array (SKA). VLBI arrays are also expected to take advantage of the increasing capacity of disk-based recording systems, and the real-time connection of antennas and correlators via commercial optical fibre networks. These upgraded and next generation telescopes will routinely reach noise levels which are at the very limit of what is now feasible with existing facilities.  Deep surveys of the microJy radio sky (e.g. Fomalont et al. 2002) suggest that the radio emission will be associated with a dominant population of moderate and high redshift galaxies that are subject to on-going, massive star formation.  In this paper I consider and contrast the capabilities of next generation and upgraded radio telescopes, in particular their ability to detect and image these faint and distant systems out to redshift 6. I also briefly comment on the prospects of detecting radio sources with the SKA before the epoch of re-ionisation. ",
        "conclusions": ""
    },
    "0206/astro-ph0206046_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} { Comparison of the old observations of Cepheids in the Small Magellanic Cloud from the Harvard data archive, with the recent OGLE and ASAS observations allows an estimate of their period changes. All of matched 557 Cepheids are still pulsating in the same mode. One of the Harvard Cepheid, HV 11289, has been tentatively matched to a star which is now apparently constant. Cepheids with ${\\rm \\log ~ P > 0.8 }$ show significant period changes, positive as well as negative. We found that for many stars these changes are significantly smaller than predicted by recent model calculations. Unfortunately, there are no models available for Cepheids with periods longer than approximatelly 80 days, while there are observed Cepheids with periods up to 210 days. }{Stars: evolution -- Cepheids -- Magellanic Clouds} ",
        "introduction": "%  Classical Cepheids are the most popular standard candles for extragalactic distance estimates. They are also useful for testing models of stellar interior and evolution.  Cepheids are massive Population I stars crossing the instability strip in the Hertzsprung--Russell diagram at the effective temperature ${\\rm \\log T_{eff} \\approx 3.8}$. Certainly most of them are undergoing core helium burning. There is also a possibility that a small fraction of observed Cepheid may be crossing the strip when the star is in the Hertzsprung gap and evolves on the thermal time scale. This crossing the instability strip is termed crossing I. The two that follows, termed crossings II and III, occur during helium burning and are generally much slower.  While a star crosses the instability strip its pulsation period changes. Even for massive objects the evolutionary period changes are very slow and a long time interval is needed to detect them. Some Cepheids in our Galaxy has been observed for about 200 years, e.g. $\\delta$ Cep, the prototype of this group, was discovered by John Goodricke in 1784.  For many stars significant period changes were detected (Berdnikov and Ignatova 2000). Recently there were also published extensive observations of some Galactic Cepheids exhibiting large period and/or amplitude changes, which are unlikely of evolutionary origin, e.g. Polaris (Kamper and Fernie 1998, Evans et al. 2002) or Y Oph (Fernie 1995). Also Turner (1998) presented data on period changes of 137 northern hemisphere Cepheids. Earlier a quantitative relation between the observed changes and those predicted by the evolutionary models was investigated by Hofmeister (1967). Saitou (1989) tried to find effects of metal abundence on the evolutionary period changes and concluded that there is a marginally dependence. However, this reasoning was based on 37 stars only and the influence of errors was not taken into account. Recently Macri, Sasselov and Stanek (2001) reported on a dramatic change in the light curve of a Cepheid discovered by Hubble in M33. They suggest that the star stopped pulsating.  In the Magellanic Clouds almost four thousand Cepheids are known. The first large database, containing periods, moments of maxima and magnitudes of the SMC Cepheids, was published by Payne-Gaposchkin and Gaposchkin (1966, hereafter PG\\&G). It is a result of long time photographic survey conducted in the Harvard observatories in the years 1888 -- 1962. Later Deasy and Wayman (1985) found that about 40 percent of a sample of 115 stars showed period variations, apparently too rapid to be explained with the evolutionary models. In the late 1990's a rich observational material  for the Magellanic Cloud Cepheids was obtained by several groups searching for gravitational microlensing. In this paper we determine period changes in the SMC Cepheids comparing the data published by PG\\&G with the results of two recent projects: OGLE (the Optical Gravitational Lensing Experiment, Udalski et al. 1997), and ASAS (the All Sky Automated Survey, Pojma\\'nski 2000). We also compare the observed period changes with the predictions of the recent stellar evolutionary models. ",
        "conclusions": "%  Our analysis led to several interesting conclusions. We identified only one Cepheid: HV 11289, which may have stopped pulsating between the Harvard and OGLE and ASAS epochs. This has to be verified especially because of possible wrong coordinates given by PG\\&G. We do not expect to find many such stars. The evolutionary models for ${\\rm Z=0.004}$ predict the time of the crossing in the loop phase as $\\sim 10^5$ for ${\\rm P=10}$ days, and $\\sim 10^4$ for ${\\rm P=30}$ days. Hence the probability of leaving the instability strip in one century by a long period Cepheid is approximately $5 \\times 10^{-3}$. The probability is smaller for Cepheids with shorter periods, i.e. less massive stars. Similar estimates are valid for the probability of mode switching. Except for HV 11289 all matched Harvard Cepheids are still pulsating in the same mode, which is not surprising in view of our estimate.  We found that no star is undergoing the first crossing, which represents a rapid evolution on the thermal time scale. Theory predicts the first crossing time to be a few tens times shorter than times for crossings II or III. Therefore, we expected to find several Cepheids with period changes corresponding to the crossing of Hertzsprung gap. Only one star, a first overtone pulsator HV 12937, has a very large rate of period change (cf. Fig. 6), but it is negative, i.e. it cannot correspond to the first crossing. This situation is similar to that presented in Paper I for the LMC sample.  We found that Cepheids with ${\\rm log ~ P > 0.8}$ have significant period changes, but for many of them one cannot decide which crossing (II or III) they are undergoing. Generally the changes are small and a few times slower than the lowest values for crossings on the nuclear time scale calculated from models given by Bono et al. (2000), as presented in Fig. 7. Moreover, we found that their calculations for ${\\rm Z=0.004}$ and ${\\rm Y=0.23}$ predict that a star with the mass ${\\rm 12 M_\\odot}$ makes a small loop but it does not reach the instability strip during the core helium burning. The discrepancy between the ABHA models and the observations for long period Cepheids is even larger. This theoretical survey also predicts too few long period Cepheids. Meanwhile, the observations confirm that there are Cepheids with periods up to 210 days in the Small Magellanic Cloud. Therefore we conclude that the predictions for massive stars, i.e. long period Cepheids, given by both sets of models cannot be right. For the first overtone Cepheids the expected and the observed rates of period changes are of the same order of magnitude, but in this case the accuracy is insufficient for a firm conclusion.  There is a good prospect for improvement in observational constraints on the rate of period changes. In near future (5--10 years) one will be able to phase them and achieve much better estimates of the period changes and their errors. Our present analysis relates only to about a half of the Cepheids listed in Harvard archives. New observations covering the entire Magellanic Clouds regions will extend the catalog. A comparison between predicted and observed period changes for hundreds of Cepheids in our Galaxy would be also interesting and would help to refine the theory of stellar structure and evolution.  \\Acknow{  I would like to thank Dr. G. Pojma\\'nski for providing photometric data on the brightest SMC Cepheids before publishing, Dr. A. Schwarzenberg-Czerny for software useful to search precise periods and Dr. P. Moskalik for providing the list of previous papers on the Cepheid period changes. I also thank Dr. S. Cassisi for making available the set of evolutionary tracks calculated by Bono et al. (2000). I am greatful to Dr. W. Dziembowski for providing the pulsation code and important discussions. I would like to thank Dr. K. Z. Stanek for useful comments and Dr. B. Paczy\\'nski for remarks and an insight to one of the original papers containing Harvard data. I wish to thank I. Soszy\\'nski and K. \\.Zebru\\'n, OGLE team members, for valuable explanations and helpful software. Support by the BW grant to Warsaw University Observatory is acknowledged. }"
    },
    "0206/astro-ph0206336_arXiv.txt": {
        "abstract": "We reconsider the scenario in which the knee in the cosmic ray spectrum is explained as due to a change in the escape mechanism of cosmic rays from the Galaxy from one dominated by transverse diffusion to one dominated by drifts. We solve the diffusion equations adopting realistic galactic field models and using diffusion coefficients appropriate for strong turbulence (with a Kolmogorov spectrum of fluctuations) and consistent with the assumed magnetic fields. We show that properly taking into account these effects leads to a natural explanation of the knee in the spectrum, and a transition towards a heavier composition above the knee is predicted. ",
        "introduction": "A well established feature of the full cosmic ray (CR) energy spectrum is that it has a power-law behavior with a steepening taking place at the so-called knee, corresponding to an energy $E_{knee}\\simeq3\\times 10^{15}$~eV. Although the knee of the CR spectrum has been known since more than four decades, none of the numerous models proposed so far to explain this feature has managed to become broadly accepted. Some proposals focus on a possible crossover between different acceleration  mechanisms below and above the knee \\cite{laga,bier,drury}, or exploit the possibility of a change in the particle acceleration efficiency \\cite{fich,joki,koba}. Other hypotheses include the nuclear photodisintegration at the sources \\cite{kara,nos}, the recent explosion of a single source \\cite{erly}, and leakage from the Galaxy due to a change in the confinement efficiency of CRs by galactic magnetic fields \\cite{syro,wdow,ptus}. Among the latter, Ptuskin et al. \\cite{ptus} in particular consider the knee as due to a crossover from a diffusive regime dominated by transverse diffusion at low energies to another dominated by drifts, i.e. by the Hall (antisymmetric) diffusion, above the knee. Since the diffusion coefficients determine the residence time of the CRs inside the Galaxy, this means that the observed slope of the CR energy spectrum will differ from the slope of the original average source spectrum just due to the energy dependence of the diffusion coefficients. Hence, the scenario just mentioned naturally accounts for a change in the spectral slope, arising from the diverse energy dependence of the two different diffusion coefficients.  In Ref.~\\cite{ptus} some important simplifying assumptions had to be done to make the differential equations determining the CR densities more tractable. In particular, a simplistic spatial dependence of the regular field and of the diffusion coefficients was adopted, but it was however found that the spectral slope beyond the knee was quite sensitive to  the magnetic field configuration adopted and to the spatial distribution of the sources. Also, the expressions for the diffusion coefficients used were actually only valid in the limit of small turbulence. Anyhow, a feature similar to the knee was obtained, but to reproduce the correct value of $E_{knee}$ a magnetic field spectrum flatter than Kolmogorov was needed. This last requirement resulted from the fact that the perpendicular diffusion coefficient (the only one relevant for the determination of the CR densities below the knee) was normalized at low energies (few~GeV) to the average value resulting from a comparison with the results of a simplified leaky box model of the Galaxy, and was then extrapolated up to $E_{knee}$ using the fact that its energy dependence is directly related to the spectral index of the random magnetic fields.   In the present work we want to further elaborate on the proposal of Ref.~\\cite{ptus}, but using more realistic configurations for the galactic magnetic fields (including e.g. field reversals, which enhance field gradients and can hence affect the CR drifts). We also  consider diffusion coefficients appropriate for strongly turbulent magnetic fields, since strong turbulence from disturbances in the interstellar plasma is thought  to be the actual picture of magnetic fields in the Galaxy, and we will directly adopt diffusion coefficients normalized from the results of numerical simulations valid in the relevant energy range of the knee and consistent with the magnetic fields adopted.  The beauty of this scenario is that in it the presence of a break in the spectrum does not require any special assumption, but just follows from properly taking into account well established properties of the propagation of charged particles in regular and turbulent magnetic fields. Moreover, this break naturally happens at the energies at which the knee is actually observed. ",
        "conclusions": ""
    },
    "0206/astro-ph0206100_arXiv.txt": {
        "abstract": "The coming GAIA Cornerstone mission by ESA will provide micro-arcsec astrometry, $\\sim$10 bands photometry and far-red spectroscopy for a huge number of stars in the Galaxy (10$^9$). GAIA spectroscopy will cover the range 8480--8740 \\AA\\, which includes the CaII triplet and the head of the Paschen series. In this paper we address the diagnostic potential of this wavelength range toward detection of peculiar stars. ",
        "introduction": "The large impact that the Hipparcos astrometric mission by ESA had on many fields of astrophysics is well known to the community. This is even more remarkable by considering that the Hipparcos observations were complete to just $V\\sim$8 mag and the horizon for astrometric errors less than 10\\% ($R_{10\\%}$) was limited to 0.1 kpc, i.e. the solar neighborhood. Nevertheless, Hipparcos data have been the main driver or at least a contributor to more than 1900 papers since the release to the community of the mission data in 1997.  Hipparcos was still flying when the European community begun openly speaking of its successor, with ideas already well in focus about GAIA by the time of the Cambridge 1995 ESA colloquium on the future of astrometry in space (ESA SP-379). Since then, GAIA has been at the center of a continent-wide effort dealing with its science goals and the technical design, culminated with the formal mission approval in the fall of 2000. Perryman et al. (2001, and references therein) provides an useful introduction to GAIA.  Before GAIA (http://astro.estec.esa.nl/GAIA/) which launch is scheduled for not later than 2012, two other survey astrometric missions could fly if their current financial problems will be eventually overcome (cf. Table~1): the German DIVA (http://www.ari.uni-heidelberg.de/diva/) and the USA mission FAME (http://www.usno.navy.mil/FAME/), both with astrometric goals intermediate between those of Hipparcos and GAIA. Finally, the technological demonstrator for interferometry in space SIM (http://sim.jpl.nasa.gov/), base-lined by NASA for a launch in 2009, will obtain micro-arcsec astrometry of a preselected limited sample of objects. While all satellites will perform photometry in parallel with astrometry, only GAIA will also collect spectra with the main aim of deriving radial velocities and therefore to measure the 6$^{th}$ component of the phase-space (the other five being provided by astrometry).   \\begin{table}[!t] \\caption{Comparison between the performances of Hipparcos, DIVA, FAME, SIM and GAIA astrometric missions (courtesy F.Mignard).} \\begin{tabular}{lccccc} &\\\\ \\tableline &Hipparcos&DIVA&FAME&SIM&GAIA\\\\ \\tableline Mission Type & scanning & scanning & scanning & pointing & scanning \\\\ Input Catalogue & Yes & No & Yes & Yes & No \\\\ V$_{completness}$ (mag) & 8 &12.5  &14  & -- & 20 \\\\ N. of objects & 118,218 & $3\\cdot 10^7$ & $4\\cdot 10^7$ & 20,000 & $1\\cdot 10^9$ \\\\ $\\sigma_\\pi$ ($\\mu$as)& 1000$_{(V=9)}$ & 200$_{(V=9)}$ & 50$_{(V=9)}$ & 4$_{(V=16)}$ & 10$_{(V=15)}$\\\\ R$_{10\\%}$ (kpc) & 0.1 & 0.5 & 2 & 25 & 10 \\\\ Spectra and RV & No & No & No & No & Yes \\\\ \\tableline\\tableline \\end{tabular} \\end{table} ",
        "conclusions": ""
    },
    "0206/astro-ph0206385_arXiv.txt": {
        "abstract": "The $Z$-burst mechanism invoked to explain ultra-high energy cosmic rays is severely constrained by measurements of the cosmic gamma-ray background by EGRET. We discuss the case of optically thick sources and show that jets and hot spots of active galaxies cannot provide the optical depth required to suppress the photon flux. Other extragalactic accelerators (AGN cores and sites of gamma ray bursts), if they are optically thick, could be tested by future measurements of the secondary neutrino flux. ",
        "introduction": "\\label{sec:intro} Recent observational data on ultra-high energy (UHE, $E\\gtrsim 10^{19}$~eV) cosmic rays give a significant evidence for clustering in their arrival directions \\cite{clustering}. This fact suggests that the observed extensive air showers are caused by particles created by point-like sources.  Recently~\\cite{BLLac} correlations of the arrival directions of cosmic rays with BL Lac type objects -- certain active galaxies located at cosmological distances -- were found (for earlier discussion see Ref.~\\cite{Farrar}). Taken seriously, these data suggest that there exist particles which can travel for cosmological distances unattenuated (without significant energy loss).  Among the Standard Model particles, only neutrinos can propagate through the Universe unattenuated at ultra-high energy. However, neutrino primaries are excluded by reconstruction of atmospheric shower development \\cite{FEnu,no-nu-primaries}. One of the ways out is to explore the so-called ``$Z$-burst'' mechanism \\cite{other-res} which works as follows. The Hot Big Bang cosmology predicts the existence of cosmic relic neutrino background.  Ultra-high energy neutrinos interact with these background neutrinos very weakly, unless the energy is fine tuned to the resonance \\cite{Zburst} with $Z$ boson production in the s-channel, that is, \\begin{equation} E_{\\rm res}\\approx {4~{\\rm eV}\\over m_\\nu}\\cdot 10^{21}~{\\rm eV}, \\label{Eres} \\end{equation} for the conventional cosmological model.\\footnote{Note that recently obtained limits on neutrino mass \\cite{Fukugita} suggest $m_\\nu\\lesssim 1$~eV.} On resonance, the interaction cross section increases significantly.  If the resonant scattering takes place within $\\sim 50$~Mpc from the Earth, then secondary protons and photons produced in decays of virtual $Z$ bosons can serve as primaries of the extensive air showers.  In astrophysical accelerators, the neutrino production is usually dominated by the two channels, namely, $p\\gamma$ and $pp$ collisions, where one proton has extremely high energy. If the collision energy in the center-of-mass frame $E_{\\rm cm}\\gg1$~GeV, the total cross section is saturated by multipion production and UHE neutrinos emerge mostly as the products of charged pion decays. For $p\\gamma$ processes the collision energy in the center-of-mass frame may be smaller, $E_{\\rm cm}\\sim1$~GeV, if UHE protons scatter off background soft photons ($E_\\gamma\\lesssim10^{-2}$~eV). In this case the cross section is saturated by production of hadronic resonances (particular type of the resonance depends on the energy in the center-of-mass frame). These resonances decay into pions, protons and neutrons.  Then UHE neutrinos appear as products of charged pion and neutron decays.  The generic feature of this mechanism is that it produces a certain amount of protons and photons per each neutrino. This may lead to contradiction with the observed fluxes if these particles leave the source.  In the case of nucleons this statement is known as the Waxman-Bahcall bound~\\cite{Waxman} (see also Ref.~\\cite{new-kalashev}): the charged cosmic ray (CR) flux above $3\\cdot 10^{18}$~eV is measured with relatively good accuracy and implies that the sources have to be opaque for UHE nucleons.  A similar situation takes place for UHE photons which escape from the source.  In the intergalactic space, the UHE photons give rise to electromagnetic cascade transferring energy into less energetic photons which propagate without attenuation \\cite{transfer-to-EGRET}.  The measurements of the flux of photons with $3\\cdot 10^7~$eV$<E_\\gamma<10^{11}$\\ eV by EGRET \\cite{egret} constrains UHE neutrino flux in a similar way as the measurement of the charged CR flux \\cite{BerezinskyS}.  Though a detailed study of the propagation should take into account a number of processes and numerical simulations are required (see, e.g., Ref.~\\cite{Semikoz}), even simple order-of-magnitude estimates demonstrate that the most part of the primary photon energy flux transfers to the EGRET energies.  A few ways were suggested which could help to overcome the EGRET limits (recent discussions on these issues can be found, for instance, in Refs.~\\cite{N7,longR,Fodor}). One option is that the sources are more abundant at distances closer to the Earth which, in average, results in less energy losses \\cite{N7}.  Alternatively, one may assume that sources are optically thick for photons as well as nucleons, and only neutrinos escape. In this paper we analyze the latter possibility. ",
        "conclusions": "\\label{sec:concl} Together with neutrinos required for the $Z$-burst explanation of UHE cosmic rays, a certain number of energetic photons are produced in astrophysical accelerators. Propagation of these photons in the intergalactic space results in their reprocessing into softer gamma rays with energies detectable by EGRET. If high-energy photons escape from the sources, the resulting gamma-ray background is inconsistent with EGRET observations for uniform source distribution. The sources, therefore, have to be opaque to ultra-high energy photons. We estimated the required mass and the size of the absorption region and found that jets and hot spots of active galaxies (blazars, radio galaxies, etc.) cannot be optically thick to UHE photons. They are, therefore, excluded as ``engines'' of the Z-burst scenario. We then analyzed the flux of the secondary neutrinos which are produced in the process of photon absorption in a hypothetical optically thick source and obtained further constraints. We argued that future measurements of neutrino flux at high energies will test the ``$Z$-burst plus optically thick protonic sources'' scenario. Our results are relevant not only for the $Z$-burst mechanism but for any other mechanism which requires high flux of photons in the source of cosmic rays."
    },
    "0206/astro-ph0206450_arXiv.txt": {
        "abstract": "We present an X-ray and radio study of the wind nebula surrounding the central pulsar PSR J1811$-$1925 in the supernova remnant G11.2$-$0.3. Using high resolution data obtained with the {\\it Chandra} X-ray observatory and with the VLA radio telescope we show the X-ray and radio emission is asymmetric around the pulsar, despite the latter's central position in the very circular shell. The new X-ray data allows us to separate the synchrotron emission of the pulsar wind nebula from the surrounding thermal emission and that from the pulsar itself. Based on X-ray data from two epochs, we observe temporal variation of the location of X-ray hot spots near the pulsar, indicating relativistic motion. We compare thermal emission observed within the shell, which may be associated with the forward shock of the pulsar wind nebula, to thermal emission from a nearby portion of the remnant shell, the temperature of which implies an expansion velocity consistent with the identification of the remnant with the historical event of 386 A.D. The measured X-ray and radio spectral indices of the nebula synchrotron emission are found to be consistent with a single synchrotron cooling break. The magnetic field implied by the break frequency is anomalously large given the apparent size and age of the nebula if a spherical morphology is assumed, but is consistent with a bipolar morphology. ",
        "introduction": "The overall structure and evolution of the X-ray emission produced by a spherically symmetric relativistic particle wind coming from a young pulsar in a supernova remnant has been modelled extensively \\citep{rc84,kc84,vand01,bcf01}. In addition, some progress has been made on the fine structure and time dependence of the sites of shock acceleration in the inner regions \\citep{ga94}. The observational tests of these models have been dominated by discussions of the Crab nebula. The advantages of studying the Crab, with its high luminosity, known age, and well observed pulsar, are many. However, the Crab is less than an ideal test case for understanding pulsar wind nebulae (PWNe) in supernova remnants (SNRs) as a class due to its anomalously high luminosity and lack of an observable SNR shell. There are a number of observed radio composite SNRs which consist of a shell and a spectrally distinct inner nebula, presumably PWNe. Only in a few cases, however,  has the central pulsar been detected (eg. W44, see Kaspi \\& Helfand 2002\\nocite{kh02} for a review, also Gotthelf 2001\\nocite{got01}), leaving the spin-down energy driving the emission to be inferred from the PWN itself.  G11.2$-$0.3 is a remarkably spherical young SNR plausibly associated with the historical event of 386 A.D. recorded by Chinese astrologers \\citep{cs77}. A 65~ms pulsar with spin-down energy $\\dot E=6.4\\times 10^{36}\\,{\\rm ergs}\\,{\\rm s}^{-1}$ was discovered in X-rays by {\\it ASCA} \\citep{tori97,tori99}. The characteristic spin-down age is much greater than the apparent age of the SNR and that inferred from its remarkably central position within the remnant given any reasonable transverse velocity (Kaspi et al. 2001\\nocite{kasp01}, hereafter Paper I). The simplest explanation for this apparent age discrepancy is that the initial spin period of the pulsar is very near the current spin period, which in turn suggests that $\\dot E$ has remained nearly constant since the supernova explosion. The distance to G11.2$-$0.3 has been reasonably well determined to be $d\\sim 5~$kpc from HI measurements \\citep{bmd85,gree88}. We therefore have a system where the assumptions of spherical symmetry, at least in regards to the shell around the pulsar,  and constant known energy input are observationally supported, and reasonable estimates of the physical size of various components can be made.  The bright spherical radio shell, reminiscent of the remnants of Tycho's and Kepler's supernovae \\citep{down84}, and the lack of an obvious hard central component in the $ROSAT$ PSPC image, initially led some authors to conclude that G11.2$-$0.3 was the remnant of a Type Ia supernova \\citep{rey94}. The first hints of a central plerionic component came from high frequency single dish radio measurements which suggested a central flattening of the radio spectrum \\citep{mr87}. {\\it ASCA} observations first clearly demonstrated that there was a central, non-thermal X-ray source \\citep{vasi96}. This component was resolved by $Chandra$, the image of which has been presented in Paper I. Further single dish radio observations have clearly demonstrated the existence of a flat central radio component \\citep{kr01}. Archival VLA observations separate the central plerion from the shell and measure the spectral index to be $\\alpha_r\\sim 0.25$ (defined as $S=S_r(\\nu/\\nu_r)^{-\\alpha_r}$, where $S_r$ is the energy flux density at frequency $\\nu_r$), a fairly typical value (Tam, Roberts, \\& Kaspi 2002,\\nocite{trk02} hereafter Paper II). In this paper we combine the radio data with the $Chandra$ data to make spectral measurements of the pulsar itself, the hard pulsar wind nebula, and soft thermal emission possibly related to the PWN which we compare to a portion of the shell. We also present evidence for apparent temporal variation in the positions of emission features in the central portion of the X-ray PWN, and make a detailed comparison of the X-ray morphology with a new 3.5~cm VLA image. We then compare these observations to models of PWN evolution and explore the possibility that within the SNR shell a wind termination shock produces the variable hard emission, and the forward shock of the PWN expanding into the stellar ejecta produces the interior thermal structures. ",
        "conclusions": "We have measured the X-ray spectrum of various components of the shell and pulsar wind nebula of G11.2$-$0.3 and compared them with the radio spectrum and morphology. The shell temperature and hard X-ray structure are consistent with a $\\sim 2000$ yr old remnant nearing the onset of the Sedov phase. The pulsar wind nebula itself shows evidence of relativistic dynamic evolution of bright X-ray enhancements near the pulsar. We argue that the reverse shock has not yet reached the expanding pulsar wind bubble, and therefore the observed structure is largely determined by the structure of the wind itself. The morphology and magnetic field suggested by the low spectral break energy suggests the outflowing wind is highly non-spherical. The wind may be largely constrained to a narrow, bipolar outflow, or possibly a thin, equatorial sheet. Further high-resolution observations at X-ray, radio and infrared wavelengths could identify the region of spot formation, help determine the true morphology, and further constrain the break energy of the spectrum."
    },
    "0206/nucl-th0206022_arXiv.txt": {
        "abstract": "A systematic study of 2-quasiparticle bands of the proton-rich odd-odd nuclei in the mass $A\\sim 70-80$ region is performed using the projected shell model approach. The study includes Br-, Rb-, and Y-isotopes with N = Z+2, and Z+4. We describe the energy spectra and electromagnetic transition strengths in terms of the configuration mixing of the angular-momentum projected multi-quasiparticle states. Signature splitting and signature inversion in the rotational bands are discussed and are shown to be well described. A preliminary study of the odd-odd N = Z nucleus, $^{74}$Rb using the concept of spontaneous symmetry breaking is also presented. ",
        "introduction": "Deformed odd-odd nuclei provide valuable information on the interplay between the collective and the single-neutron and -proton degrees of freedom in atomic nuclei. In general, odd-odd nuclei are difficult to study both experimentally and theoretically due to the complexity of their low-lying spectra. In comparison to even-even nuclei where the pairing correlations favor a particular configuration, in odd-odd nuclei many configurations are equally probable, and therefore, contribute to the complexity of the spectra. A detailed experimental analysis for the odd-odd nuclei with mass $A\\sim 70-80$ has become possible only in the last few years with the availability of large arrays of high resolution HPGe detectors. Some extensive measurements of the proton-rich odd-odd nuclei \\cite {70br1,70br2,72br1,74br1,74br2,74rb1,76rb1,78rb1,78rb2,78y1,80y1,82y1} have been carried out for the Br-, Rb- and Y-isotopes in laboratories around the world. The excited states of some of these nuclei have been populated upto high-spins, e.g. $I \\approx 20 \\hbar$, by heavy-ion fusion reactions. In each of these nuclei several bands have been established. The low-lying levels of these bands show very irregular behavior, whereas at moderate rotational-frequency, i.e. $\\hbar\\omega \\approx 0.5$ MeV, regular rotational bands start developing. The stretched $E2$ transition strengths deduced from measured lifetimes of excited states indicate higher degree collectivity for the high-spin states. Furthermore, a considerable amount of odd-even staggering (signature splitting) in transition energies is observed in these bands. In some of them, the staggering phase shows an inversion  at spin $I\\approx 10$.  The study of these proton-rich nuclei is not only interesting from the nuclear structure point of view, but also has important implications in nuclear astrophysics \\cite{rp}. Since heavier elements are made in stellar evolution and  explosions, nuclear physics, and in particular nuclear structure far from stability, enters into the stellar modeling in a crucial way. It is believed that these proton-rich nuclei near the $N=Z$ line are synthesised in the rapid-proton capture process (the rp-process) under appropriate astrophysical conditions. The X-ray burst is suggested as a possible site. Recently, the proton-rich odd-odd nuclei with $A\\sim 70-80$ have been studied more rigorously partially due to this astrophysical interest. Understanding the structure of the low-lying states in these nuclei is also important for the study of Fermi superallowed $\\beta^+$ decays \\cite{rp3}.  On the theoretical front, very little effort has been put in the study of doubly odd nuclei as compared to even-even nuclei in this mass region. In an early work, the rotational band structure of doubly odd nuclei in this mass region has been investigated in $^{76}$Br \\cite {kr79} on the basis of 2-quasiparticle rotor model. More recently, the signature inversion phenomenon in these nuclei have been studied in detail via the axially symmetric rotor plus quasiparticle model \\cite {re01}. The microscopic calculations based on the Excited-VAMPIR approach \\cite{pa1,pa2} were performed for some odd-odd nuclei. The large-scale spherical shell model \\cite{cau94} which is quite successful for describing the $pf$-shell nuclei, can not be applied for these well-deformed mass-80 nuclei since it is important to include the $g-$shell. The configuration space required for such a study is quite enormous and can not be handled by the present state-of-the-art computational facilties.  In recent years, the projected shell model (PSM) \\cite{ha95} has become quite successful in explaining a broad range of properties of deformed nuclei in various regions of the nuclear periodic-table. The most striking aspect of this quantum mechanical model is its ability to describe the finer details of the high-spin spectroscopy data with simple physical interpretations. The studies of  odd-odd rare-earth nuclei \\cite{ha91} have shown this capability. Very recently, the PSM has been systematically applied to the even-even Kr-, Sr-, and Zr-isotopes of the proton-rich region\\cite {pa01}, where it has been found that the PSM results are able to explain most of the experimental observations. The purpose of the present work is to carry out a similar study for the proton-rich, odd-odd Br-, Rb- and Y-isotopes with N = Z+2 and Z+4, and with an example of N = Z nucleus. The physical quantities to be described are energy spectrum and electromagnetic transition probability. A quantitative comparison of calculated and measured transition strengths can provide a stringent test of the model. We would like to mention here that the present study is the first major application of the projected shell model approach to the medium mass odd-odd nuclei. It is also for the first time that calculations for the electromagnetic transition strengths in odd-odd nuclei are performed within the PSM framework.  The manuscript is organized as follows: In the next section, we shall give a brief outline of the model for completeness. This section also gives expressions for various physical quantities to be discussed in this paper. The results of calculations and comparisons with the experimental data are presented in Section 3. Finally, a summary is given in Section 4. ",
        "conclusions": "For N = Z+2, Z+4 isotopes of odd-odd Br, Rb, and Y, the $\\pi g_{9/2} \\otimes \\nu g_{9/2}$, 2-qp configuration has been firmly established for the yrast region. This configuration has been observed up to or above 10 MeV excitation energy in these nuclei. Therefore, we first discuss the calculations for these nuclei in the following three subsections. The discussion includes the band diagrams, transition energies and electromagnetic transition strengths (the B(M1) and B(E2) values) of each of the above mentioned isotopes.  The odd-odd N = Z nuclei of this mass region appear to have a different structure as compared to the neighbouring nuclei. Experimentally, the odd-odd N = Z nuclei are far more difficult to study since they lie further away from the stability valley, approaching the proton drip-line. For the few measured cases \\cite{70br1,70br2,74rb1}, both the low-spin and high-spin states exhibit distinct properties. The structure with the presence of the isospin, T = 1 and T = 0 bands can be further complicated by the prolate-oblate shape co-existence at low spin for many nuclei. Therefore, we preferred to choose the experimentally known N = Z nucleus $^{74}$Rb \\cite{74rb1} which has large prolate deformation even in ground state. Separate description of T = 1 and T = 0 bands was given to show the applicability of the simple model based on schematic Hamiltonian to this extreme case. However, the present version of PSM does not treat explicitly the isospin degree of freedom to calculate the interaction between T= 1 and T = 0 bands, giving a constraint for the actual descriptin of the structure.   The PSM calculations proceed in two steps. First, an optimum set of deformed basis is constructed from the standard Nilsson model. The Nilsson parameters are taken from Ref. \\cite {su00} and the calculations are performed by considering three major shells (N = 2, 3 and 4) for both neutrons and protons. This basis is large enough so that all nucleons in the f-p-g shells are active. The basis deformation $\\epsilon_2$ used for each nucleus is given in Table I. These values are taken either according to experimental information, if available, or from theoretical calculations. We emphasize that unlike the cranking mean field approaches, the deformation parameters used as an input to the PSM calculations need not correspond exactly to the true nuclear deformation. This is because of the shell-model nature of the PSM: The deformed single-particle states serve solely as a way to truncate shell-model basis. All observable properties in the PSM calculations are determined by the many-body wave-functions obtained by diagonalising the shell model Hamiltonian. In principle, the larger the deviation of the initial deformation from the true one, the bigger will be the configuration space for an actual description of the nuclei. In the present calculations, the space is truncated by the inclusion of the states within an energy window of 3.5 MeV around the Fermi surface. This determines the size of the basis space, $ | \\phi_\\kappa \\ket $ in Eq. (\\ref{intrinsic}), of the order of 30. In the second step, these basis states are projected to good angular momentum states, and the projected basis is then used to diagonalize the shell model Hamiltonian. The diagonalization gives rise to the energy spectra, and the transition strengths are subsequently calculated using the resulting wave functions.   \\subsection{Band diagram}  For the deformation values given in Table I, the $[440]\\frac {1}{2},~[431]\\frac{3}{2},~[422]\\frac{5}{2}$ orbitals for protons and neutrons lie near the Fermi surfaces in the nuclei under consideration. To describe the low-lying structure of these nuclei, the 2-qp states based on these orbitals (with $K = K_\\nu \\pm K_\\pi$) should be taken into account. A representative band diagram for $^{74}$Br, containing some unperturbed rotational bands, is shown in Fig. 1 to see the underlying structure. Note that only the important configurations are displayed in the band diagram, although many more others are included in the calculations.  Out of the all unperturbed bands taken in the diagonalization, only few of them show energy staggering between odd- and even-spin states (the so-called signature splitting). As can be seen in Fig. 1, these bands have the zigzag behavior as a function of spin. This occurs for the bands with configurations involving low-K orbitals. It was noticed in Ref. \\cite {ha91} that unlike in the particle rotor model, the configurations based on orbitals with $K \\neq \\frac{1}{2}$ also contribute to the energy staggering. The unperturbed bands which show staggering influence the yrast states through the band mixing, and thus cause the signature dependence in the yrast band (shown as filled circles in Fig. 1). It is seen from Fig. 1 that the lowest $K=4$ band has only small staggering. The clear signature splitting at around spin $I=10$ come mainly from the mixing of the 2-qp band based on $( \\nu [431]\\frac{3}{2} \\otimes \\pi  [431]\\frac{3}{2} )$. Two more pairs of bands with K = 1, 2 based on $( \\nu [431]\\frac{3}{2} \\otimes \\pi [440]\\frac{1}{2} )$ and K = 1, 0 based on $( \\nu [440]\\frac{1}{2} \\otimes \\pi [440]\\frac{1}{2} )$ have much stronger signature splitting. They dive down to the yrast region at higher spins, and therefore bring strong energy staggering to the yrast band. It is thus the configuration mixing that produces nicely the observed signature splitting as one will see in Fig. 2 below.  Interestingly, the current PSM calculations indicate that the phase of the energy staggering may change at certain spins. In Fig. 1, two of the above mentioned bands with K = 0 based on $( \\nu [440]\\frac{1}{2} \\otimes \\pi [440]\\frac{1}{2} )$ and K = 1 based on $( \\nu [431]\\frac{3}{2} \\otimes \\pi [440]\\frac{1}{2} )$ show an inversion in the staggering phase (the so-called signature inversion). The inversion appears at $I=6$ for the former, and $I=10$ for the latter. Note that this is a signature inversion occurring within one unperturbed band, without involving any dynamic mixture with other bands. Hara and Sun called this kind of inversion a self-inversion, to distinguish this from an inversion involving two bands with mutually opposite staggering phases \\cite {ha91}. The phenomenon of self-inversion predicted in Ref. \\cite {ha91} has not been confirmed experimentally. Here, the band with self-inversion (emphasized by a solid bold curve at $I=8-11$ in Fig. 1) lies at very low excitations and the inversion appears at spin $I\\approx 10$. This is a promising candidate that may explain the observed signature inversion in the rotational band in $^{74}$Br \\cite{74br1} and other nuclei in this mass region.  \\subsection{Transition energies}  After diagonalization at each spin, the lowest energies give the energies of the yrast states. The energy differences between adjacent spin states are compared with the measured transition energies in Fig. 2 for $^{72,74}$Br, $^{76,78}$Rb and $^{80,82}$Y. The characteristic feature of the six nuclei is the exhibition of a clear energy staggering between the odd- and even-spin states, as seen in the $E(I)-E(I-1)$ plot in Fig. 2. It is called signature splitting because the odd- and even-spin states can be classified as two groups specified by the signature quantum number. In addition, an inversion in the staggering phase is observed in most of the nuclei. Signature inversion in the mass-80 region was discussed in Ref. \\cite{re01}.  As can be seen from Fig. 2, for all of the nuclei the agreement between the calculation and experiment is quite satisfactory above $I \\approx 10 \\hbar$. The energy splitting at higher spins is well reproduced in all the cases, indicating that the important influence on the yrast band from the low-K components of the $g_{9/2}$ valance neutrons and protons are correctly accounted for by the configuration mixing.  However, the calculation does not reproduce the signature inversion observed at low spins in $^{74}$Br, $^{76,78}$Rb and $^{80,82}$Y. This is rather unfortunate because the mechanism that may cause the inversion is clearly present in the model. In this medium mass region with $N\\sim Z$, rotational bands with mutually opposite signature dependence can not co-exist near the yrast region, and therefore, the only possibility to explain the observed signature inversion is through the self-inversion. In Ref. \\cite{ha91}, the favorite condition for this to happen was given to nuclei where bands based on configurations with K = 0 $( \\nu K=\\frac{1}{2} \\otimes \\pi K=\\frac{1}{2} )$ and K = 1  $( \\nu K=\\frac{3}{2} \\otimes \\pi K=\\frac{1}{2} )$ lie low in energy. $^{74}$Br fulfils these conditions. The failure that our final results do not reproduce the signature inversion may indicate that the band mixing is not strong enough to bring the desired feature of self-inversion into the yrast band. One may generally argue that there should be an interaction that particularly pushes the unperturbed band with self-inversion further down to affect the yrast states more strongly. Obviously, this force is absent in the schematic Hamiltonian employed in the current model.  \\subsection{Electromagnetic transition strengths}  The experimental data from lifetime measurements performed for $^{74}$Br, $^{78}$Rb and $^{82}$Y \\cite {74br1,74br2,78rb1,78rb2,82y1} have separately provided the B(M1) and B(E2) values. For $^{72}$Br and $^{76}$Rb, no lifetime measurements have been performed: the ratio B(M1)/B(E2) could be obtained from the measured intensity of transitions. In the present section, these data are compared with our calculations.  The experimental B(M1) values for $^{74}$Br, $^{78}$Rb and $^{82}$Y are plotted in Fig 3. In all the three nuclei, staggering between the odd- and even-spin states is clearly present. This variation in the B(M1) pattern is directly compared with the calculated transition strengths. Except for a few particular states, the calculation achieves a quantitatively agreement with data. In contrast to the transition energies, no inversion in the staggering phase is obtained in both theory and experiment. The measured B(E2) values for $^{74}$Br, $^{78}$Rb and $^{82}$Y are plotted in Fig. 4. The global trend of the B(E2) curve as a function of spin is qualitatively reproduced by the calculation as shown in the Fig. 4, although some deviations are seen at high spins.  In Fig. 5, the calculated B(M1)/B(E2) values are compared with the experimental data wherever available for $^{72}$Br, $^{76}$Rb and $^{80}$Y. A staggering pattern is again observed in experimentally deduced B(M1)/B(E2) ratios for those levels for which cascade M1 transitions have been observed. Indeed, in $^{72}$Br B(M1)/B(E2) values for the 11$^{(+)}$ and 13$^{(+)}$ levels are $\\approx~10~ \\mu_n^2/(eb)^2 $, while for 10$^{(+)}$ level its value is only 0.03(1) $\\mu_n^2/(eb)^2 $ \\cite {72br1}. These features are quite nicely reproduced by the calculation. The experimental ratio B(M1)/B(E2) in $^{76}$Rb were determined from the measured intensity of transitions reported in Ref. \\cite {76rb1}. The B(M1)/B(E2) values in $^{76}$Rb are generally smaller compared to the ratio observed in $^{72}$Br. This fact reflects the smaller contribution from the lower K orbitals for $^{76}$Rb. Unfortunately, in case of $^{80}$Y no intensities for different transitions in the positive parity band is available to allow a determination of the B(M1)/B(E2) ratios experimentally. The calculated values of B(M1)/B(E2) ratio are small compared to those of $^{72}$Br and $^{76}$Rb.  \\subsection{The N = Z nucleus $^{74}$Rb}  In Ref. \\cite{74rb1}, two rotational bands were observed in the N = Z nucleus $^{74}$Rb. In contrast to the above discussed N = Z+2 and Z+4 nuclei, the ground-state band of $^{74}$Rb does not look like a rotational band having the usual 2-qp structure. Instead, the transition energies of the ground state band in $^{74}$Rb show similarities to those in even-even nuclei. In fact, they have been interpreted \\cite{74rb1} as being formed from the T = 1 isobaric analog states of $^{74}$Kr with pairing correlations based on T = 1 neutron-proton pairs. At higher rotational frequency, a T = 0 rotational band becomes energetically favored over the T = 1 ground state band.  As a full description for these observations is not possible in the present version of the PSM since the physics is beyond what the model space in (\\ref{intrinsic}) can provide, we tried to describe separately T = 0 and T = 1 bands. Using the concept of spontaneous breaking of the isospin invariance, in ref. \\cite{FS99} it was argued that the neutron-proton interaction can be effectively considered in a theory with only neutron-neutron and proton-proton pairings, since one has a freedom to choose a one particular direction in the isospin space. One thus ends up with the understanding that the T = 1 ground state band in an odd-odd N = Z  nucleus may be approached by a 0-qp state with nn and pp pairings and the T = 0 rotational band has a structure similar to the 2-qp states \\cite{FS99}.  The two neighboring even-even N = Z nuclei $^{72}$Kr and $^{76}$Sr have a large deformation with $\\epsilon_2 \\approx 0.36$ \\cite{pa01,SS01}. It is thus reasonable to perform a calculation for $^{74}$Rb with this deformation. At the deformation $\\epsilon_2 = 0.36$, the [422]$\\frac{5}{2}$ orbital is the closest to the respective neutron and proton Fermi levels in $^{74}$Rb. Thus, the band based on the 2-qp state of $K=5$ ($\\nu$[422]$\\frac{5}{2} \\otimes \\pi$[422]$\\frac{5}{2}$) is the lowest in energy among all the 2-qp states (see Fig. 6). At higher spins, bands consist of smaller $K$-states $[440]\\frac{1}{2},~ [431]\\frac{3}{2}$, which show stronger signature splitting, mix with the $K=5$ band through configuration mixing. Therefore, the mixed band follows the signature phase that favors the odd-spin states. This explains the absence of even spin states in the experimental spectra at the high spin region \\cite{74rb1}.  In Fig. 7, the two observed rotational bands in $^{74}$Rb are compared with the calculations in an $E(I) - E(I-2)$ plot. It is seen that the calculated results fit the data reasonably well. This indicates that the physical understanding of ref. \\cite{FS99} on the band structure in an odd-odd N=Z nucleus is correct. However, in order to understand the interplay between the T=0 and 1 bands, it is necessary to include the neutron-proton pairing in the shell model Hamiltonian."
    },
    "0206/astro-ph0206251_arXiv.txt": {
        "abstract": "We present a case study of validating an astrophysical simulation code. Our study focuses on validating FLASH, a parallel, adaptive-mesh hydrodynamics code for studying the compressible, reactive flows found in many astrophysical environments. We describe the astrophysics problems of interest and the challenges associated with simulating these problems. We describe methodology and discuss solutions to difficulties encountered in verification and validation. We describe verification tests regularly administered to the code, present the results of new verification tests, and outline a method for testing general equations of state. We present the results of two validation tests in which we compared simulations to experimental data. The first is of a laser-driven shock propagating through a multi-layer target, a configuration subject to both Rayleigh-Taylor and Richtmyer-Meshkov instabilities. The second test is a classic Rayleigh-Taylor instability, where a heavy fluid is supported against the force of gravity by a light fluid. Our simulations of the multi-layer target experiments showed good agreement with the experimental results, but our simulations of the Rayleigh-Taylor instability did not agree well with the experimental results. We discuss our findings and present results of additional simulations undertaken to further investigate the Rayleigh-Taylor instability. ",
        "introduction": "The enormous progress seen in the evolution of fast computing machines and numerical methods stands as one of the great achievements of the twentieth century. Numerical modeling is now an accepted and widely applied method of research, and in many cases simulations have matured to the extent that they now provide direction to theoretical research. In astrophysics, where the complexity of many of the problems requires large-scale computing for any hope of progress, much of research involves the development and application of reliable and trustworthy simulation codes. Before the scientific community can have confidence in results from such a code, it must be subjected to a wide variety of verification and validation tests. This paper discusses progress in verification and validation of FLASH, a parallel, adaptive-mesh simulation code for the compressible, reactive flows found in many astrophysical environments.  The goal motivating the development of FLASH is to advance the solution of several astrophysical problems related to thermonuclear flashes on the surfaces and in the interiors of compact objects. In particular, the problems of interest are type I X-ray bursts, classical novae, and Type Ia supernovae. These events all involve the accretion of material from a companion star onto the surface of the compact star, followed by the ignition of either the core of the compact star or the material accreted onto the surface. The global physical phenomena common to all three of these events include an accretion flow onto the surfaces of compact stars, shear flow and Rayleigh-Taylor instabilities~\\citep{taylor50,chandra81} on the stellar surfaces and in the core, ignition of thermonuclear burning in degenerate matter, development of convection, propagation of nuclear burning fronts, and expansion of the stellar envelope. An understanding of these global phenomena requires knowledge of the fundamental physical processes involved in each. Accordingly, much of our scientific effort focuses on research into the basic ``microphysics.\" These fundamental processes include turbulence at large Reynolds and Rayleigh numbers, fluid instabilities and mixing, convection and the convective penetration of stable matter at very high densities, thermodynamics in relativistic and degenerate regimes, the propagation of both subsonic and supersonic burning fronts, and radiation hydrodynamics.  Verification and validation are fundamental steps in developing any new technology, whether it be a simulation code like FLASH or an instrument for observation. For simulation technology, the goal of these testing steps is assessing the credibility of modeling and simulation. Considerable work on verification and validation of simulations has been done in the field of computational fluid dynamics (CFD), and in the CFD literature the terms verification and validation have precise, technical meanings~\\citep{aiaa98,roache98a,roache98b}. Verification is taken to mean demonstrating that a code or simulation accurately represents the conceptual model. Validation of a simulation means demonstrating that the simulation appropriately describes nature. The scope of validation is therefore much larger than that of verification and includes comparison of numerical results with experimental or observational data. In astrophysics, where it is difficult to obtain observations suitable for comparison to numerical simulations, this process can present unique challenges.  In this paper, we describe our efforts at verifying and validating the hydrodynamics module in FLASH. We begin by describing the astrophysical problems of interest and the importance of fluid instabilities in the problems. A discussion of the terminology and methodology for verification and validation (V\\&V) follows this description. We follow that with a discussion of the challenges found in verifying and validating astrophysical simulations, and in particular, the aspects of astrophysical modeling for which it is difficult to apply recommended CFD V\\&V techniques. In the next section, we present verification tests and include an outline of a procedure for testing an arbitrary equation of state in the context of numerical hydrodynamics schemes. The following section contains a comparison of the results from two laboratory experiments with those obtained from simulations, and the final section contains discussion and conclusions.  \\subsection{Overview of Astrophysical Thermonuclear Flashes}  Thermonuclear flashes, events of rapid or explosive thermonuclear burning, occur in a variety of stellar settings. These events include type I X-ray bursts, classical novae, and Type Ia supernovae, all of which involve a close binary system in which matter from a companion star accretes onto the surface of a compact star (neutron star or white dwarf).  Either the core of the compact object or the accreted layer on the surface of the compact object ignites under electron-degenerate conditions, and a thermonuclear burning front is born and begins to propagate.  These events provide not only fantastic observational displays, but also tools with which potentially to answer several fundamental questions. The light curves and spectra of X-ray bursts can provide information about the masses and radii of neutron stars~\\citep{lewin93,lamb01} and thus also provide information about the nuclear equation of state. Classical novae can provide information about the abundances of intermediate-mass elements in the universe and the dynamics of white dwarfs in close binary systems~\\citep{gehrz98}. Type Ia supernovae provide additional information about the abundances of intermediate-mass and heavy elements and play a crucial role as ``standard candles\" in determining cosmological parameters such as the Hubble constant, $H_0$, the mass density, $\\Omega_M$, and the cosmological constant or vacuum energy density, $\\Omega_\\Lambda$~\\citep[see][and references therein]{riess98,perlmutter98,turner01}.  X-ray bursts are flashes that start at the bottom of a very thin layer ($\\sim 10-100$ m) of hydrogen-rich or helium-rich fuel that has accreted onto the surface of a neutron star \\citep{taam85,lewin93,taam93}. The total energy released by burning the fuel into ash is a factor of $\\sim 20-100$ less than the gravitational binding energy. Consequently, the accreted material is gravitationally bound to the neutron star and the flash is not quenched by expansion of the envelope~\\citep{hansen75}.  Instead, fuel in the accreted envelope is incinerated to iron-peak or heavier nuclei~\\citep{shatz01}.  Novae result from the ignition of a layer ($\\sim 10^4$ m) of hydrogen-rich material that has accreted onto the surface of a white dwarf \\citep{truran82,shara89,starrfield89,livio94}.  In this case, the total energy released by thermonuclear burning is a factor of $\\sim$ 100 more than the gravitational binding energy.  As a result, the excess energy leads to an enormous expansion of the white dwarf's envelope, which engulfs the companion star and forms a common envelope binary. The work done against gravity in the expansion of the envelope cools the hydrogen burning layer and quenches the runaway, leading to the establishment of a phase of stable hydrogen burning that continues through envelope exhaustion.  Type Ia supernovae are thought to be due to carbon flashes that ignite in the cores of accreting white dwarfs \\citep{woosley86, nomoto94, niemeyer95a, niemeyer95b}.  Models involving either a pure deflagration or a pure detonation have been unable to provide a consistent explanation for the observed expansion velocities and the spectrum of intermediate-mass and iron-peak ejecta. Some Type Ia supernova models that involve a transition from a deflagration to a detonation have been constructed.  One possibility is a more-or-less spontaneous transformation of the initial deflagration into a detonation as the burning front propagates outward through the white dwarf \\citep{niemeyer95a,niemeyer95b,khokhlov97,niemeyer97}. Another possibility is that the initial deflagration dies out as a result of the expansion of the outer layers of the white dwarf. When these gravitationally bound layers collapse back onto the white dwarf, a detonation is ignited \\citep{blinnikov87, boisseau96, khokhlov95, khokhlov97}.  In either case, these models are capable of accounting for the observed expansion velocities of the silicon-group and iron-group nuclei.  In all three of these thermonuclear flash events, the nuclear burning time scale is much shorter than the time scale over which the nuclear fuel accretes onto the surface. These short burning time scales make it likely that ignition of the fuel occurs at either a single point or, at most, at a few discrete points.  The situation may be complicated by the presence of magnetic fields.  The strong magnetic fields (B $\\sim 10^6$ - $10^9$ G) of white dwarfs and the super-strong magnetic fields (B $\\sim 10^9$ - $10^{12}$ G) of neutron stars may be capable of funneling the flow of accreting matter onto the magnetic polar caps of the compact object \\citep{lamb73}. The accreted matter, which constitutes the nuclear fuel, may or may not be able to spread over the surface of the star before ignition occurs.  The three-dimensional nature of this fuel geometry is compounded by the effects of a magnetic field on the thermal and mass transport coefficients \\citep{lamb96,potekhin99a,potekhin99b}.  Even in the absence of a magnetic field, thermonuclear flashes are inherently multi-dimensional because of the complexity of the underlying fluid instabilities. Most efforts at modeling the multi-dimensional nature of thermonuclear flashes have been two-dimensional \\citep{fryxell82a, steinmetz92, shankar92, shankar94, livne93, glasner95, glasner97, khokhlov95, boisseau96, kercek98}, but there have been a few three-dimensional studies \\citep{khokhlov95, khokhlov97, garcia99, kercek99}.  Recent progress includes a detailed two-dimensional study of helium detonations on the surface of a neutron star that confirmed that a detonation can spread burning over the entire surface on a time scale consistent with burst rise times~\\citep{zingale01}, three-dimensional simulations of thermonuclear explosions of Chandrasekhar-mass carbon-oxygen white dwarfs~\\citep{reinecke99}, and studies of the cellular structure of detonation fronts in both two and three dimensions~\\citep{timmes00a,timmes01}.  \\subsection{The Role of Fluid Instabilities in Flash Problems}  Fluid instabilities and subsequent mixing are expected to play a fundamental role in the events involving thermonuclear flashes.  For example, determining whether or not there is substantial mixing between the accreted hydrogen-helium envelope and the carbon-oxygen surface layer of the white dwarf is crucial to understanding the nova mechanism. Mixing of intermediate-mass elements into the accreted layer is critical because otherwise hydrogen burning would be too slow to produce a nova. Furthermore, without this mixing it is difficult to produce the observed abundances of intermediate-mass nuclei in the ejecta~\\citep[see][and references therein]{rosner01}.  Such mixing may occur during the first phases of the accretion cycle or when convection in the accreted layer works its way down to the interface with the white dwarf. In the latter case, convective undershoot may dredge up material from the white dwarf and mix it into the accreted layer \\citep{livio91}. Although some progress has been made on modeling this nova mixing mechanism \\citep{glasner97,kercek98,kercek99}, a consensus has not been reached. Another proposed mixing mechanism, which is the subject of ongoing research, is breaking of nonlinear resonant gravity waves at the carbon-oxygen surface \\citep{rosner00,rosner01,alexakis01,alexakis02}.  In the case of a neutron star, penetration into the crust by convection in the accreted layer is strongly inhibited by the large jump in atomic weight between the heavy element crust and the hydrogen-helium composition of the accreted layer.  Thus, no significant mixing by this sort of convective dredge-up is expected to occur in thermonuclear flashes involving the surface layers of a neutron star. Some mixing may occur via other mechanisms such as shear instabilities, but, unlike the nova case, the mixing of the inert heavy elements from the underlying neutron star into the accreted layer is not likely to have a significant effect on the evolution.  Fluid instabilities and mixing are also expected to play a key role in the explosion mechanism of a Type Ia supernova.  A subsonic burning front that begins near the center of a massive white dwarf is subject to Kelvin-Helmholtz, Landau-Darrieus, and Rayleigh-Taylor instabilities \\citep{khokhlov97, khokhlov01, hillebrandt01}.  Growth of these instabilities dramatically increases the surface area of the burning front. This increase in surface area increases both the burning rate and the speed of the front. The dependence of the speed of the burning front on fluid instabilities is one of the reasons a study of Rayleigh-Taylor instabilities is a key component in our efforts at V\\&V.  Because fluid instabilities play a fundamental role in thermonuclear flash events and  may be probed in reasonably good terrestrial experiments, experiments involving fluid instabilities have been the focus of our validation efforts thus far. In what follows, after describing the methodology of V\\&V and briefly describing our numerical methods,  we present verification tests and the results of two validation problems-- a laser-driven shock propagating through a multi-layer target and the classic Rayleigh-Taylor problem.   \\subsection{The Process of Verification and Validation}  The testing step of code development, which involves verification and validation of the numerical methods and resulting simulations, is one of the most important procedures required for successful numerical modeling. This fact has been frequently overlooked in astrophysics. In our description of V\\&V, we take the point of view that the code to be tested is either one under development or an existing code that is being applied to a new problem. The details of testing the code will apply in either case. The results of testing will feed back to the choices of numerical methods for addressing the physics. If a particular method or model fails a test, then another must be chosen. We note that verification and validation are necessary but not sufficient tests for determining whether a code is working properly or a modeling effort is successful. These tests can only determine for certain that a code is not working properly.  V\\&V is a maturing area of study in the field of CFD, and there is a wealth of information available in the literature~\\citep[cf.][]{aiaa98,oberkampf98,pilch01,roache98a,roache98b}. The fundamental strategy of V\\&V is the assessment of error and uncertainty in a computational simulation. The requisite methodology is complex because it must address sources of error in theory, experiment, and computation. Because these areas of study present diverse perspectives (and because V\\&V is still a developing field), it is common to find disagreement in the terminology of V\\&V~\\citep{aiaa98}.  We adopt the following definitions from the American Institute of Aeronautics and Astronautics (AIAA)~\\citep{aiaa98}: \\begin{quotation} \\begin{trivlist} \\item Model: A representation of a physical system or process intended to enhance our ability to understand, predict, or control its behavior. \\item Modeling: The process of construction or modification of a model. \\item Simulation: The exercise or use of a model. (That is, a model is used in a simulation) \\item Verification: The process of determining that a model implementation accurately represents the developer's conceptual description of the model and the solution of the model. \\item Validation: The process of determining the degree to which a model is an accurate representation of the real world from the perspective of the intended uses of the model. \\item Uncertainty: A potential deficiency in any phase or activity of the modeling process that is due to lack of knowledge. \\item Error: A recognizable deficiency in any phase or activity of modeling that is not due to lack of knowledge. \\item Prediction: Use of a CFD model to foretell the state of a physical system under conditions for which the CFD model has not been validated. \\item Calibration: The process of adjusting numerical or physical modeling parameters in the computational model for the purpose of improving agreement with experimental data. \\end{trivlist} \\end{quotation} \\citet{roache98a} offers a concise, if informal, summary: \\begin{quotation} First and foremost, we must repeat the essential distinction between Code Verification and Validation. Following \\citet{boehm81} and \\citet{blottner90}, we adopt the succinct description of ``Verification\" as ``solving the equations right\", and ``Validation\" as ``solving the right equations\". The code author defines precisely what partial differential equations are being solved, and convincingly demonstrates that they are solved correctly, i.e. usually with some order of accuracy, and always consistently, so that as some measure of discretization (e.g. the mesh increments) $\\Delta \\rightarrow 0$, the code produces a solution to the continuum equations; this is Verification. Whether or not those equations and that solution bear any relation to a physical problem of interest to the code user is the subject of Validation. \\end{quotation} Roache goes on to point out that in a meaningful sense, a ``code\" cannot be validated, but only a calculation or range of calculations can be validated. He also makes the distinction between verifying a code and verifying a calculation, noting that ``use of a verified code is not enough.\" In our discussion, we will adhere to these definitions as closely as possible.  Another term requiring discussion is ``convergence.\" In CFD, the term is used in two different ways. ``Grid convergence\" refers to the convergence of the discretization error of the numerical solution as the grid size and time step approach zero. ``Iterative convergence\" refers to the convergence of the results of successive steps of an iterative procedure within a numerical method. \\citet{roache98a} notes that ``inadequate iterative convergence will pollute grid convergence results\" and that the issue of iterative convergence can blur the distinction between verifying a code and verifying a calculation because iterative tuning parameters can be problem dependent.  Figure~\\ref{fig:valscheme}, a Venn diagram~\\citep{venn80} illustrating the parts of numerical modeling, provides a schematic for considering the role of verification and validation in numerical modeling. On the left, the largest circle represents Nature, or at least the part of Nature in which the problem of interest resides. The smaller circles inside of Nature represent, from largest to smallest, the range of desired validity of the code, the range of actual validity, and the range of the design goals of the code. The circle in the center of the diagram (to the right of the Nature circle) represents the theory describing Nature (which may be thought of as the model), and the circle to the right represents numerical scheme(s) implementing the theory (modeling and simulation).  Verification begins by identifying the purported design goal of the code or code modules, which is represented by the smallest circle in the center of the Nature circle on the schematic. This is the step Roache refers to as defining ``precisely what partial differential equations are being solved,\" though the process may require a larger scope. In the language of the AIAA, this may be thought of as identifying and understanding the implementation of the model. Verification, confirming that simulations produced by the code accurately implement the design goal of the code (i.e. the implementation accurately represents the model and the solution to the model), may be thought of as confirming the ``mapping\" between the theory and numerical method circles on the schematic. The process requires identification and a quantitative description of the error, and the strategy is typically a systematic study of mesh and time step refinement as is appropriate for finite difference, finite volume, and finite element methods. Other numerical methods such as vortex, lattice gas, and Monte Carlo methods require different procedures~\\citep{aiaa98}.  Verification tests are typically simple problems, often with known analytic solutions, that can be used to study the accuracy and convergence rates of a code. Despite the relative simplicity of the tests, verification is a difficult process. Analytic solutions, which are easiest to compare against, are few and usually are limited tests of the physics. Even in these simple cases, singularities and discontinuities complicate the process of verification. Singularities arising from the geometry or the coordinate system should be removed where possible. Singularities inherent in the conceptual model (e.g. shocks and discontinuities in the flow) require special attention. These structures will typically exhibit a lower order of convergence from the rest of the flow, if indeed they lead to a converged solution at all. An added level of complexity arises in the case of discontinuous flows because not all numerical methods for solving a given set of PDEs are equivalent. Different methods handle flow discontinuities differently and/or rely on different (intrinsic numerical or physically-motivated) subgrid models. Therefore, confirming that a numerical method is correctly solving a set of PDEs can also be a validation problem.  Test problems, even without analytic solutions, may be devised to monitor conservation, symmetry properties, and effects of boundary conditions~\\citep{aiaa98}. Additional progress may be made by testing self-convergence. Self-convergence can be difficult to test, however, and changing the resolution by a only factor of two does not demonstrate a converged answer~\\citep{fryxell94}. Further, resolving the length and time scales relevant to the physical problem well enough for convergence may be prohibitively expensive. Also, simulations with complex physics, particularly three-dimensional simulations, may not have resolved solutions~\\citep{aiaa98}. We further note that convergence tests say nothing about the correctness of the answer in that it is possible for solutions to converge to the wrong answer.  Verification testing also may include code-to-code comparisons, that is, the comparison of the results of simulations performed with different codes. For these types of comparisons, care should be taken to find accurate, benchmarked solutions calculated very carefully by independent investigators, preferably using different numerical approaches~\\citep{aiaa98}. Examples in astrophysics and cosmology of code-to-code comparisons include the Geospace Environmental Modeling (GEM) Magnetic Reconnection Challenge~\\citep{birn01} and the Santa Barbara Cluster Comparison Project~\\citep{frenk99}. The GEM Reconnection Challenge is a collaborative study of phenomenon of magnetic reconnection, which plays an important role in magnetosphere dynamics. The Santa Barbara Cluster Comparison Project compared the results of twelve codes simulating the formation of a galaxy cluster in a flat cold dark matter universe.  Both studies performed simulations starting from a uniform set of initial conditions, with the goal of studying how different numerical methods capture the behavior of the system under study.  Diagnostics were constructed that were sensitive to the physical processes essential to these systems and robust enough to compare across codes.  The utility of code-to-code comparisons both for finding bugs and for increasing the understanding of the behavior of different numerical methods in complex situations is clear. This procedure can only be meaningful, however, if the codes involved have undergone other rigorous V\\&V tests.  As with code-to-analytic or code-to-experimental comparisons, these tests can only indicate possible problems. Code-to-code comparisons can also include consistency tests such as the nightly comparison tests that FLASH undergoes, which will be discussed below.  Validating a simulation requires identifying the key elements involved in the simulations and for each element (as well as the integrated code) constructing test problems that have the results of laboratory experiments as the accepted results. Key elements include parts of the code that describe the fundamental physical processes and include items such as transitions to low and high Mach numbers, transport of energy by conduction or radiation, transport of energy by advection (convection), source terms (e.g. nuclear burning), equations of state, opacities, and microscopic transport such as molecular diffusion and viscosity. Validation problems tend to be much more complex than verification problems and typically have no analytic solutions. These problems nevertheless involve phenomena that are sufficiently simple to be studied both by experiment and simulation, and these problems are used to determine whether a particular calculation reproduces the outcome of the phenomenon or experiment.  We note that validation goes beyond purely numerical testing; it includes testing the fundamental assumptions and concepts that go into a model and probing the range of validity of a model. In the schematic, validation may be thought of as probing the area of the actual validity circle. An example that has implications for the problems here is that of compressibility.  As \\citet{roache98a} mentions, a thoroughly verified incompressible fluid dynamics code will produce invalid results when applied to a problem in which compressibility of the fluid affects the dynamics. The situation is more complicated, however, if one considers the validity of applying a compressible code to an incompressible problem (or a very low Mach number flow). In this case, the compressible code may correctly address the problem, but the CFL limit will require a huge number of very small time steps, making the problem intractable and the solution less accurate due to the accumulation of numerical errors. Thus incompressible (or very low Mach number) flow may be within the formal range of validity of a compressible code (i.e., the range that is mathematically well-posed), but not within the actual or practical range of validity.  The challenges associated with validating an astrophysical simulation code exceed those of validating a standard fluid dynamics code. Astrophysical events are often complex and involve many interacting physical processes, each of which must be tested. Validating any simulation code with laboratory experiments can be a difficult process, particularly if the diagnostic resolution of the experiment is poor, if there are significant uncertainties in the material properties, or if the initial/boundary conditions are not well-defined. The situation is especially acute in astrophysics, where we are limited to observations of distant objects. The interiors of stars, for instance, do not lend themselves to direct observation, and, even if this were not a problem, the vast majority of astrophysical objects are too far away to resolve. Observations of thermonuclear flashes can only show the results of the events (light curves and spectra) and not the details of initiation of the outburst. Also, the length scales of the astrophysical objects present challenges. Flows within stars, for example, are expected to occur at Reynolds numbers greater than 10$^9$, and terrestrial experiments cannot approach such a regime.  In addition, validation testing includes systematic grid sensitivity studies to asses grid convergence error and assess the level of refinement necessary to capture the key physical effects~\\citep{aiaa98}. Accordingly, many of the difficulties in astrophysical verification also appear in validation. As noted, the flows of interest in astrophysics, particularly flows in thermonuclear flash events, involve unsteady flows, and convergence testing is more difficult for these flows. The required accuracy of validation activities, however, is not generally as stringent as that of verification activities~\\citep{aiaa98}.  Another issue making astrophysical validation more complicated than CFD validation is the choice of equation of state. The CFD literature we studied mentions equations of state as possible sources of error, but does not emphasize the importance of choice of equation of state on the validity of the model. Phenomena such as degeneracy and other interactions play a critical role in astrophysics, and the influence of these enter the simulation through the equation of state. The properties of an equation of state (e.g. convexity) are particularly important when using hydrodynamics methods that solve the Riemann problem as does the principal hydrodynamics module in FLASH~\\citep{menikoff89}. Considerable effort can go into validating a hydrodynamics code with a known, accurate equation of state, but coupling it to an inappropriate or inaccurate equation of state will invalidate any simulations. Also, effort expended in testing equations of state relevant for laboratory experiments will not necessarily improve confidence in predictions made by astrophysical simulations that require another equation of state.  Because of these complexities, it seems unlikely that one could ever satisfactorily validate an astrophysical simulation. Instead, validation efforts focus on laboratory experiments that capture the relevant physics, with the expectation that the experience gained from these closely related cases builds confidence in the predictions of the astrophysical simulations.  Accordingly, a significant part of the challenge of validating astrophysical simulation codes is to find acceptable laboratory experiments.  A good experiment for validation should be a good experiment itself (that is, it should provide accepted, repeatable results), it should adequately capture a significant portion of the physical processes of interest, and it should be diagnosed well enough for a meaningful comparison to simulation.  A final subject to describe in the methodology of V\\&V is calibration. Calibration is not validation. Instead, calibration is a process performed in order to improve the agreement of computational results with experiments, and calibration does not generate the same level of predictive confidence as validation. Calibration is performed when there is uncertainty in the modeling of complex processes and also when there are incomplete or imprecise measurements in the experiments. Calibration involves adjustments to parameters in subgrid models, reaction rates, and boundary conditions, and it includes assumptions about minimal or optimal levels of mesh refinement that are made in cases where there are not completely resolved solutions~\\citep{aiaa98}.  We complete our introduction with the observation that much of what we have said about simulation validation is true of validating any theoretical model of a system, including more traditional analytic models. ",
        "conclusions": "In this paper, we presented the results of our efforts at verification and validation of FLASH, our astrophysical simulation code. We presented the results of new verification tests and the results of two validation tests in which we were able to carefully compare our simulation results to experimental results. These tests served to increase our understanding of the physics relevant to the problems of interest, to improve our simulation techniques, and to build confidence in our results.  \\subsection{Verification} The verification tests we performed for this study indicate that the PPM hydrodynamics module is performing as expected. Advection tests on a uniform mesh showed slightly better than second order convergence of the error with mesh resolution.  The PPM module, on a uniform mesh, propagated a sound wave for one period with approximately second order accuracy, and a shock wave, as expected, to first order. In the case of a shock wave on an adaptive mesh, the code showed the expected first-order convergence as well. The code was able to hold a standing shock steady for several times steps, also, indicating the Riemann solver is working correctly.  The material equation of state, as well as the hydrodynamic solver, must be well known and well behaved in order to obtain a meaningful comparison with experiments.  Verification testing of a test equation of state resulted in establishing a procedure for testing consistency and quantifying the results. This procedure will allow us to quickly test new equations of state as we develop additional physics modules and address new astrophysical problems.  \\subsection{Validation}  The results of simulating laser-driven shock experiments show that we can capture the bulk properties of the flow, including the morphological properties of the resolvable structures. We observed the expected instabilities at the two material interfaces that grew with time during the course of the simulations. We performed a mesh resolution study that in general showed convergence of instability growth at the material interfaces. The resolution study did not completely demonstrate convergence, though, because the amount of small scale structure present in the simulations increased with resolution as may be seen by a visual inspection of the results (Figure~\\ref{fig:6ch}). We attribute this increasing amount of small scale structure to the fact that the amount of numerical dissipation in our hydrodynamics method decreases with increasing resolution. Visual inspection of the experimental results and our simulated radiographs suggested that a simulation with six levels of adaptive mesh refinement better matched the observed small scale structure. We note, though, that the correct amount of small scale structure in the experiments is not likely captured in the experimental diagnostics because of resolution limits and because the radiograph produces a two-dimensional shadowgraph of a three-dimensional experiment, possibly averaging out small structures.  Measurement of the lengths of the spikes allowed us to quantify the results of the simulations, and we found that simulations with seven and eight levels of adaptive mesh refinement agreed to within 5\\% of the highest resolution (nine levels) simulation at all times during the evolution (Figure~\\ref{fig:7andexp4}).   Complete convergence is prohibited by the growth in small-scale structure with increasing resolution.  The spike lengths at intermediate resolutions match the experimental results very well, falling within the experimental errors.  We interpret these as an important calibration result that indicates that six and seven levels of adaptive mesh refinement are appropriate for two-dimensional simulations of these particular experiments.  In addition, the simulations showed that the shock propagating through the corrugated material interface develops a perturbation that oscillates and becomes planar at the second interface, as expected from theory. We also observed the expected imprinting of the perturbation on this second interface.  For this validation problem, comparison between simulation and experiment is limited by the diagnostic resolution of the experiments and the amount of physics included in the simulations.  Improvements to experimental diagnostics will be made in the next generation of laser experiments, and improvements to the simulations could be made by inclusion of a more realistic material equation of state, adding the walls of the shock tube, and modeling the laser-driven energy deposition process.  Absent these improvements, we cannot conclude that the good agreement of our results with the experiment completely validates the simulations.  The results of simulations of a single-mode Rayleigh-Taylor instability demonstrated the effect of resolution and dimensionality on growth rates. The single-mode simulations established that $\\sim$ 25 grid points per mode are required for a reasonable estimate for the growth rate. This calibration result suggests a minimum resolution for simulations involving the Rayleigh-Taylor instability with our code. This required resolution may make complete simulations of many problems prohibitively expensive at the present time.  We confirm previous findings \\citep{kane00,young01} that instability growth rates in three-dimensional simulations are larger than those found in the equivalent two-dimensional simulations, which indicates limits on results from two-dimensional models of astrophysical phenomena.  The single-mode simulations also showed that, as with the laser-driven shock simulations, one sees structure formation at smaller and smaller scales as resolution is increased. The result is that increasing resolution will not necessarily produce a converged flow. Further, the single-mode simulations also showed that small scale structure in highly resolved simulations can have an effect on integral properties of the flow, in this case instability growth rates.  Our multi-mode Rayleigh-Taylor instability simulations, like the simulations of the laser experiment, show that the code results agree with the observed bulk properties of the flow: we observe a mixing zone that is very similar in structure to that of the experiment.  Our simulations do not, however, agree well with the experimental rate coefficient, $\\alpha$; our results are systematically lower.  This difference could be due to limitations of either the simulations or the experiments. The experimental initial conditions, which strongly effect the growth rate, are poorly characterized.  Further, limitations on experimental diagnostics may distort measured growth rates and mixing.  The simulations, on the other hand, may have been under-resolved.  In our low-resolution simulation, most of the power in the initial conditions was in modes which were resolved with only $2$--$4$ grid points; in light of our single-mode results, this is far too low. Increasing the resolution to resolve these modes with $4$--$8$ points increases the observed growth rate, in (somewhat) closer agreement with experimental results.  \\subsection{Evaluation of V\\&V Results}  The principal conclusions we draw from our efforts are that validation of an astrophysical simulation code is a difficult process and that the process led us to unanticipated questions. The verification tests we performed allowed for a careful quantitative study of the accuracy and convergence rates of FLASH. The exercise proved useful in a variety of ways, not the least of which was developing techniques for accurately comparing the results of our simulations to accepted answers. The validation tests we performed, while they did not conclusively validate the simulations, provided insight into the many issues involved in numerical modeling, and served to calibrate the code and build confidence in our results.  V\\&V testing allows for an assessment of error in the code and progress in building confidence in the results. The process is limited, though. Analytic solutions are typically available only for simple problems, which may not serve as a strenuous tests of a particular code. Convergence studies, though essential, do not imply that the converged answer is correct. Code-to-code comparisons are another useful tool; they can be more probing than analytic solutions, and can shed insight into the behavior of different methods. But again, similar answers do not imply correctness.  Comparing numerical and experimental results is also difficult. Nonlinear systems typically have exponential sensitivity to initial conditions, so that any unmodeled initial perturbation in the system can greatly affect the results.  Experimental results may be limited, noisy measurements of a poorly-characterized but real physical situations, while the numerical results may be complete diagnostics of the situation derived by a model of dubious applicability.  Comparisons must, therefore, rely on choosing features which are obtainable from both results, yet sensitive enough to the physics of interest that comparing the features is a strong probe of the model and numerics.  The complexity of this procedure means that doing it well requires feedback between the computational and experimental researchers, with all parties improving techniques and performing multiple runs.  The computational simulations may be very expensive, however, and performing relevant experiments is a costly and difficult process.   This makes it hard to improve results and demonstrate reproducibility.  As our stated goal in this work was validating an astrophysical simulation code, the assumptions in the models and their applicability to both terrestrial and astrophysical flows warrants discussion. Our numerical models did not have an explicit viscosity. We saw, however, that the effects of numerical dissipation cannot be avoided in simulations.  Other physical processes that may play a role but are not included in the models include surface tension, species diffusion, and thermal diffusion. With only Euler's equations, it is impossible to adjust the relevant dimensionless numbers (Schmidt, Reynolds, Rayleigh, and Prandtl numbers, for example) to match the correct properties of the fluid.  Our use of models without an explicit viscosity resulted in an increasing amount of small scale structure observed with increasing resolution in simulations of both validation problems. As mentioned above, limits on diagnostic resolution prevent us from determining the correct amount of small scale structure in both validation experiments. Were the experimental diagnostics better, we could incorporate physically-motivated terms for viscosity, surface tension, or diffusion into the equations we evolve and calibrate their magnitudes until we obtain a converged flow with the correct amount of small scale structure. Proceeding in this manner (as well as including a physically-motivated equation of state) would help us validate our simulations of these and similar terrestrial experiments, but would be of little use in the astrophysical case, where viscosity and surface tension play no dynamic role.  Even with better diagnostics these types of experiments can only further validate the code if we model additional physics that are not relevant to the astrophysical problems of interest. As we mentioned above, flows within stars are expected to have a Reynolds number greater than 10$^9$, for which the viscosity-free Euler equations are a better approximation, but such flows are impossible to achieve either in the laboratory or in current simulations. Therefore, we  conclude that the terrestrial experiments we have simulated have served to build confidence in the hydrodynamics module in FLASH, but that there are certainly still limits on the strength of statements we can make about the validity of FLASH simulations of laboratory experiments.  We end by repeating a point from our introduction, namely that verification and validation can determine only if a code returns an incorrect answer. Verification and validation cannot prove that a theoretical result, either numeric or analytic, is correct. By doing a sufficient number of tests, however, one can significantly increase one's confidence in the results. Our efforts at validation, although they have presented many challenges and led to new questions such as the effect of small scale structure on bulk properties of flows, have increased our confidence in the simulations produced by FLASH."
    },
    "0206/astro-ph0206121_arXiv.txt": {
        "abstract": "We present a high resolution ($\\lambda/\\Delta\\lambda\\approx 500$) X-ray spectrum of the bright quasar H1821+643 ($z=0.297$), obtained in a 470 ksec observation with the \\chandra X-ray Observatory. We search for X-ray absorption by highly ionized metal species, \\ovii and \\oviii in particular, at the redshifts of the six intervening \\ovi absorption systems known from UV studies.  We detect features with $\\ga 2\\sigma$ significance at the predicted \\ovii and \\oviii wavelengths of one \\ovi system, at the \\ovii wavelength of a second, and at the \\neix wavelength of a third. We find two additional features of comparable strength (one \\ovii and one \\oviiin) within $1000\\;\\kms$ of \\ovi redshifts.  The $1\\sigma$ constraints on the relative abundances of different species imply significant variations from system to system in $f(\\ovin)$, the fraction of oxygen in the \\ovi state. The constraints in the two detected \\ovi systems imply gas overdensities lower than the values $\\delta \\ga 100$ expected in virialized systems, suggesting that the absorption arises in lower density, filamentary structures. At the $2\\sigma$ level, however, the physical constraints are weak, though all of the systems must have temperature $T < 10^6\\K$ to be consistent with upper limits on \\oviin. If we treat our $2\\sigma$ detections of known \\ovi systems as real, but assume minimal \\ovii and \\oviii in the other systems, we estimate $[f(\\ovin)+f(\\oviin)+f(\\oviiin)]/f(\\ovin)=32 \\pm 9$ for the average ratio of all highly ionized oxygen species to \\ovin.  Combined with estimates of the total column density of \\ovi absorption per unit redshift, this ratio implies that the total baryon fraction associated with detected \\ovi absorbers is $\\Omega_b(\\ovin)\\sim 0.03 h_{70}^{-1}$, a substantial fraction of the baryon density predicted by big bang nucleosynthesis, and larger than that associated with stars or with gas detected in 21cm or X-ray emission. Because of the limited S/N of the detections, these results must be treated with caution.  Nonetheless, the combination of the \\ovi data with these X-ray forest measurements provides the most direct evidence to date for the pervasive, moderate density, shock-heated intergalactic medium predicted by leading cosmological scenarios.  The high inferred incidence of relatively strong \\ovii and \\oviii absorption implies that some regions of this medium are enriched to a level substantially above $[{\\rm O/H}]=-1$. ",
        "introduction": "\\label{sec:intro}  Much of our knowledge of the intergalactic medium (IGM) comes from the rest-frame UV line absorption that it imprints on the spectra of background quasars: the \\lya forest of neutral hydrogen and associated metal lines such as \\civ and \\ovi. The \\lya forest is produced mainly by diffuse, photoionized gas at temperatures $T\\sim 10^4\\K$, and it appears to trace the main reservoir of cosmic baryons at high redshift \\citep{rauch95,hernquist96,rauch97,weinberg97}. The \\lya forest thins out at low redshift, and cosmological simulations predict that the continuing process of structure formation heats a substantial fraction of intergalactic gas to temperatures where it produces little hydrogen \\lya absorption, and where the dominant ionization stages of heavier elements have absorption transitions at X-ray wavelengths rather than UV.  Hot, dense gas in the central regions of galaxy clusters and groups can be detected by its X-ray emission, but most of the shock-heated gas resides in the outskirts of virialized halos and in lower density, filamentary structures, making its continuum emission weak. One of the few prospects for detecting this low density, shock-heated gas is via the ``X-ray forest'' of absorption lines it should produce in quasar spectra \\citep{hellsten98,perna98,fang00,chen02,fang02a}. \\cite{nicastro02} find strong evidence for absorption by highly ionized oxygen and neon at $z\\approx 0$, but there have been no clear detections to date of intergalactic X-ray absorption originating beyond the Local Group. This paper describes a search for X-ray forest absorption towards H1821+643 ($z=0.297$) using a high resolution ($\\lambda/\\Delta\\lambda \\approx 500$) spectrum obtained in a 470 ksec observation with the \\chandra X-ray Observatory.  The baryon density implied by big bang nucleosynthesis and the estimated primordial deuterium abundance, $\\Omega_{\\rm BBN} \\approx 0.04 h_{70}^{-2}$ (\\citealt{burles97,burles98}; here $h_{70} \\equiv H_0/70\\;\\hubunits$), exceeds the density of baryons in known stars and X-ray emitting gas by roughly an order of magnitude \\citep{fukugita98}. The lower density regime of the ``warm-hot intergalactic medium'' (WHIM, a term coined by \\citeauthor{cen99} [\\citeyear{cen99}] to refer to gas in the temperature range $10^5-10^7\\K$) could constitute a major fraction of the ``missing'' low redshift baryons.  Hydrodynamic cosmological simulations predict that 30-50\\% of the baryons reside in this phase at $z=0$ (\\citealt{cen99}; \\citealt{dave99}, \\citeyear{dave01}). HST and FUSE detections of \\ovi ($\\lambda\\lambda 1032, 1038$\\AA) absorption lines towards H1821+643 and PG0953+415 (\\citealt{tripp00}, hereafter TSJ; \\citealt{ts00,oegerle00,savage02}) offer a tantalizing hint of this baryon reservoir. Adopting conservative assumptions of [O/H]$=-1$ and an \\ovi ionization fraction $f$(\\ovin)=0.2 (which is close to the maximum in photo- or collisional ionization), TSJ conclude that the gas associated with these weak \\ovi absorbers accounts for $\\Omega_b \\approx 0.004 h_{70}^{-1}$ of the cosmic baryon budget, comparable in total mass to all other known low redshift components combined. They remark, moreover, that $f(\\ovin)$ could plausibly be much lower than 0.2, implying a substantially higher baryon fraction. For most reasonable assumptions about the physical conditions of this absorbing gas, the dominant ionization state should be \\ovii or \\oviiin. Gas hotter than $T\\sim 5\\times 10^5\\K$ would have very low $f(\\ovin)$, so it might produce \\ovii or \\oviii absorption with no detectable \\ovin.  We chose H1821+643 for our X-ray forest search in part because it is one of the brightest X-ray quasar at moderate redshift, and in part because it has been studied carefully for intervening \\ovi absorption (TSJ; \\citealt{oegerle00}).  Theoretical expectations for X-ray forest absorption depend strongly on the assumed metallicity distribution of the moderate density, shock-heated IGM, a subject on which there is little empirical guidance.  However, for metallicities $Z \\sim 0.1-0.3 Z_\\odot$, simulations and analytic calculations predict few \\ovii or \\oviii lines above column density $N\\sim 2\\times 10^{15}\\cm^{-2}$ \\citep{hellsten98,perna98,fang00,chen02,fang02a}, and weaker absorption by other elements, so even long exposures with \\chandra or {\\it XMM-Newton} are unlikely to detect absorption at high signal-to-noise ratio.  Searching at redshifts with known \\ovi absorption allows one to adopt a lower effective threshold for significant detection, greatly increasing the odds of success. Furthermore, detections of or upper limits on X-ray absorption provide new constraints on the physical conditions of the \\ovi absorbers.  We chose the duration of our spectroscopic observation so that if $[f(\\ovin)+f(\\oviin)+f(\\oviiin)]/f(\\ovin) \\sim 30$, a physically plausible ratio that would imply that \\ovi absorbers contain a large fraction of the ``missing'' low redshift baryons, then we would clearly detect the strongest of the TSJ systems, and we would detect the next two strongest systems at the $2-3\\sigma$ level.  As we will show below, we find only a $\\sim 1\\sigma$ signal from the strongest TSJ system, but we obtain $\\approx 2\\sigma$ detections of \\ovii from the next strongest system and of \\ovii and \\oviii from the \\ovi system discovered by \\cite{oegerle00}.  We also find features that could correspond to oxygen absorption at other redshifts close to the detected systems (within 1000$\\kms$), and a possible signature of \\neix absorption from one system. The best estimate of the mean $[f(\\ovin)+f(\\oviin)+f(\\oviiin)]/f(\\ovin)$ from our data is $\\sim 30$, but the moderate statistical significance of our detections and the implied variation in $f(\\ovin)$ from system to system leave the interpretation of our results more ambiguous than we would like. Nonetheless, this observation provides the strongest evidence to date for the moderate density, shock-heated IGM predicted by theoretical models, and the detections and upper limits have important implications for the physical state of the \\ovi absorption systems and for their connection to the missing baryons of the low redshift universe.  We proceed to a description of the observations, data reduction, and absorption line analysis, then to a discussion of these implications. ",
        "conclusions": "\\label{sec:conclusions}  In combination with the {\\it HST} and {\\it FUSE} observations of \\ovi absorption, this {\\it Chandra} spectrum of H1821+643 provides the most direct evidence to date for the pervasive, moderate density, shock-heated IGM predicted by leading cosmological models.  The spectrum shows absorption signals with $\\ga 2\\sigma$ significance at the expected wavelengths of \\ovii and \\oviii absorption for one of the \\ovi systems, at the expected wavelength of \\ovii absorption for another, and at the expected wavelength of \\neix absorption for a third.  There are two absorption features of comparable strength within $1000\\kms$ of predicted \\ovii and \\oviii absorption wavelengths, which could represent additional absorbing gas in the large scale environment of the \\ovi systems.  Some of the other $\\sim 2\\sigma$ features in the spectrum could correspond to X-ray forest systems without associated \\ovi absorption, but these features are present in roughly the number expected for Gaussian noise, so without an {\\it a priori} reason to search for absorption at their observed wavelengths, we have no convincing evidence that they represent physical systems.  There is a clear detection of Galactic \\oi absorption, and a strong, broad feature at 25.35\\AA\\, for which we have no obvious identification.  The best estimates of the ratios $N(\\oviin)/N(\\ovin)$ and $N(\\oviiin)/N(\\ovin)$ for the two detected \\ovi systems imply that they are largely photoionized and that gas densities are significantly below those expected in virialized structures like groups or clusters. This inference relies on the assumption that the \\ovin, \\oviin, and \\oviii absorption arise in the same gas, not in different components of a multi-phase system.  Independent of the single-phase assumption, the measured ratios imply that the column density-weighted mean value of $[f(\\ovin)]^{-1}$ is substantially higher than the value of 5.0 that TSJ used to derive their lower limits on $\\Omega_b(\\ovin)$, the baryon density associated with \\ovi absorbers.  TSJ and \\cite{savage02} find $\\Omega_b(\\ovin) \\sim 0.002-0.004h_{70}^{-1}$ for $[{\\rm O/H}]=-1$ and $[f(\\ovin)]^{-1}=5$, and our results imply that $\\langle [f(\\ovin)]^{-1} \\rangle = 32 \\pm 9$ at $1\\sigma$.  At face value, therefore, our measurements suggest that the baryon fraction associated with \\ovi absorbers is substantially larger than that in stars, cold gas, or hot gas detected in X-ray emission, and that the \\ovi absorbers trace a significant reservoir of the ``missing'' low redshift baryons.  The candidate \\ovii and \\oviii systems that are not at the \\ovi redshifts contain a similar total amount of highly ionized oxygen.  In combination with the results of hydrodynamic simulations \\citep{hellsten98,fang02b,chen02}, the existence of even one or two X-ray forest lines at the column densities of our detected systems suggests that some regions of the shock-heated IGM have been enriched to well above $[{\\rm O/H}]=-1$, and the oxygen and hydrogen column densities measured for the z2 \\ovi absorber suggest $[{\\rm O/H}]\\ga -0.3$ in this system.  Unfortunately, our results also show that definitive measurements of X-ray forest absorption are extraordinarily difficult, even with the great technological advance that {\\it Chandra}, LETG, and ACIS-S represent over previous instruments.  The absorption signals for the detected \\ovi systems differ from zero at only the $2-3\\sigma$ level, and since we detect some but not all of the known \\ovi systems, we cannot rule out the possibility that the apparent association between these features and predicted absorption redshifts is simply a coincidence.  The features at other wavelengths are not strong enough to represent clear detections of new systems.  The constraints on the physical properties and associated baryon fraction of the \\ovi absorbers are very interesting at the $1\\sigma$ level, but with $2\\sigma$ error bars we obtain only loose constraints that are not particularly surprising, though they are still more than was previously known.  {\\it XMM-Newton} has larger effective area than {\\it Chandra} at soft X-ray wavelengths.  However, the effective area near 25\\AA\\ is non-uniform, with many instrumental features at observationally interesting wavelengths.  Furthermore, published spectra obtained with RGS at these wavelengths seem to achieve lower resolution than {\\it Chandra}/LETG, reducing the sensitivity to narrow lines.  It is not clear, therefore, whether {\\it XMM-Newton} will prove as powerful as \\chandra in X-ray forest searches.  Studies of X-ray forest absorption are among the major scientific drivers for {\\it Constellation-X},\\footnote{http://constellation.gsfc.nasa.gov} and for the still more powerful {\\it XEUS} mission.\\footnote{http://astro.estec.esa.nl/SA-general/Projects/XEUS} {\\it Chandra} and {\\it XMM-Newton} can at best give us a glimpse of what these future missions may achieve. However, the tantalizing but still ambiguous hints of X-ray forest absorption in the H1821+643 spectrum would be frustrating to live with for half a decade or more.  In our view, the best step forward in the study of the X-ray forest would be further {\\it Chandra} observations of H1821+643, which remains the best target for such an investigation because of its X-ray brightness and well studied \\ovi absorption, and because the {\\it Chandra} spectrum presented here provides 500 ksec of existing data and clear objectives for a future observation. For example, an additional 1 Msec observation would reduce the noise in the co-added spectrum by $3^{1/2}$, so any true physical features that have $2\\sigma$ significance in the current spectrum would rise to $3.5\\sigma$ significance, making their reality entirely unambiguous, regardless of whether they lie at the redshift of a known \\ovi system.  Conversely, since we obtain $1\\sigma$ upper limits $N(\\oviin) \\la 1.9\\times 10^{15}\\cm^{-2}$ and $N(\\oviiin) \\la 3\\times 10^{15}\\cm^{-2}$ for systems with no absorption signal in the current spectrum, the $1\\sigma$ upper limits in undetected systems from a co-added 1.5 Msec spectrum would imply $f(\\oviin)/f(\\ovin) < 11\\; [N(\\ovin)/10^{14}\\cm^{-2}]^{-1}$ and $f(\\oviiin)/f(\\ovin) < 18\\; [N(\\ovin)/10^{14}\\cm^{-2}]^{-1}$, respectively. These constraints are still not extremely strong for any individual system, but null detections for all of the known \\ovi systems in such a spectrum would suggest that the total amount of material they contain is not enough to represent the main reservoir of the missing low redshift baryons.  {\\it Chandra} observations can also guide the design of X-ray forest studies with future missions.  If the apparent detections in our current spectrum represent real systems, then {\\it Constellation-X} is likely to find an abundance of X-ray absorption lines tracing the enriched, shock-heated IGM.  Conversely, if the evidence for these features weakens as the S/N improves, then tall trees in the X-ray forest are rare, and even {\\it Constellation-X} will have to search long and hard to find them."
    },
    "0206/astro-ph0206192_arXiv.txt": {
        "abstract": "We present the results of the first detailed millimeter and submillimeter molecular line survey of bright rimmed clouds, observed at FCRAO in the CO~(\\jequals{1}{0}), \\ceio~(\\jequals{1}{0}), \\hcop~(\\jequals{1}{0}), \\hthcop~(\\jequals{1}{0}), and \\nthp~(\\jequals{1}{0}) transitions, and at the HHT in the CO~(\\jequals{2}{1}), \\hcop~(\\jequals{3}{2}), \\hcop~(\\jequals{4}{3}), \\hthcop~(\\jequals{3}{2}), and \\hthcop~(\\jequals{4}{3}) molecular line transitions. The source list is composed of a selection of bright rimmed clouds from the catalog of such objects compiled by \\citet{sfo}. We also present observations of three Bok globules done for comparison with the bright rimmed clouds. We find that the appearance of the millimeter CO and \\hcop\\ emission is dominated by the morphology of the shock front in the bright rimmed clouds. The \\hcop~(\\jequals{1}{0}) emission tends to trace the swept up gas ridge and overdense regions which may be triggered to collapse as a result of sequential star formation. Five of the seven bright rimmed clouds we observe seem to have an outflow, however only one shows the spectral line blue-asymmetric signature that is indicative of infall, in the optically thick \\hcop\\ emission. We also present evidence that in bright rimmed clouds the nearby shock front may heat the core from outside-in thereby washing out the normally observed line infall signatures seen in isolated star forming regions. We find that the derived core masses of these bright rimmed clouds are similar to other low and intermediate mass star forming regions. ",
        "introduction": "Triggered Star Formation is the process by which star formation is initiated or accelerated through compression of a clump in a molecular cloud by a shock front \\citep{elm1}. Although a basic understanding of the process of triggered star formation exists, only recently have numerical models becoming sophisticated enough to yield detailed comparisons with observations. Comparison of hydrodynamic shock induced star formation models to observations of potential triggered star formation regions is complicated by the fact that such regions can be very chaotic. Shocks, outflows, and core rotation all have kinematic signatures which may overwhelm the signatures associated with the triggering mechanism \\citep{elm2}. In order to alleviate this confusion somewhat, we choose to observe regions of potential small scale triggering. These regions typically involve one core in a cometary cloud, and a well defined shock front geometry.  Several numerical studies of shock triggered collapse have been undertaken recently \\citep[hereafter VC]{bos,fb1,fb2,vc}. These studies have found that shocks with velocities less than 45 km s$^{-1}$ can cause cores to collapse, but fast shocks ($> 100$ km s$^{-1}$) tend to destroy the principal cooling molecules, leading ultimately to the destruction of the cloud \\citetext{VC}. These models also yield an appreciable number of binary and multiple star systems, which are common in our galaxy, yet difficult to form by spontaneous, isolated star formation. The VC models tend to reproduce characteristics of observed cometary clouds thought to be small scale triggered star formation regions \\citep{elm1,elm2}. This comparison between the models and the observations will be explored further in a forthcoming paper.  In order to perform a detailed study of sites of triggered star formation, which can be compared to models, we conducted a survey of bright rimmed clouds identified by \\citet[hereafter, SFO]{sfo}. These are all molecular clouds which border expanding HII regions. At the boundary between the ionized gas within the Stromgren sphere, and the molecular cloud is an ionization front, which shows up as the bright rim identified by SFO. Embedded near the rim, and just within the molecular cloud in each of these sources is an IRAS source. The fact that the ionization front is collecting gas into a ridge or core in many of these sources makes them excellent candidates for the globule squeezing or collect and collapse mode of triggered star formation \\citep{elm1,elm2}. These clouds are also divided into three morphological types, shown schematically in Figure~\\ref{schematic}. The first is type~A, which has a moderately curved rim, and looks much like a shield in three dimensions. The second is type~B, which has a more tightly curved rim near the head of the cloud, but which tends to broaden near the tail. Type~B is also known as an elephant trunk morphology. The third is type~C which has a very tightly curved rim and a well defined tail. Type~C is often called a cometary cloud. The shock induced collapse models of VC as well as observations of the expanding ionization front of the Orion OB 1 association \\citep{os2} suggest that these morphological types of bright rimmed clouds may actually be a time evolution sequence with clouds evolving from type~A through type~B to type~C. ",
        "conclusions": "Our observations constitute the first detailed millimeter and submillimeter multitransition study of bright rimmed clouds. Among the 7 bright rimmed clouds we observe, 6 seem to share traits similar with other low to intermediate mass star forming regions. Our analysis of these bright rimmed clouds has yielded the principal results that follow. \\begin{enumerate} \\item New FCRAO CO~(\\jequals{1}{0}), \\ceio~(\\jequals{1}{0}), \\hcop~(\\jequals{1}{0}), \\hthcop~(\\jequals{1}{0}), and \\nthp~(\\jequals{1}{0}) observations along with new HHT CO~(\\jequals{2}{1}), \\hcop~(\\jequals{3}{2}), \\hcop~(\\jequals{4}{3}), \\hthcop~(\\jequals{3}{2}), and \\hthcop~(\\jequals{4}{3}) observations of 7 bright rimmed clouds and 3 Bok globules were presented. These observations constitute the most detailed millimeter and submillimeter study of bright rimmed clouds to date. \\item The millimeter CO and \\hcop\\ emission tends to terminate abruptly at the ionization front. As a result, the overall morphology of the CO and \\hcop\\ millimeter integrated intensity maps are similar with the optical morphologies identified by SFO. \\item The millimeter \\hcop\\ tends to show the dense swept up ridge behind the ionization front, as well as the star forming core around the embedded IRAS source. In some of the bright rimmed clouds the \\hcop~(\\jequals{1}{0}) emission also traces other overdense clumps which may later be triggered to collapse  by the ionization front, resulting in sequential star formation. \\item The millimeter and submillimeter \\hcop\\ lines from many of the bright rimmed clouds appear nearly gaussian, with little evidence of infall asymmetry. The only exceptions to this are SFO~18, which shows significant blue asymmetry, and SFO~16 which shows a slight red asymmetry relative to optically thin tracers. \\item The core masses derived for the bright rimmed clouds using both \\nthp\\ and \\hcop\\ are typical for low and intermediate mass star formation regions. The \\nthp\\ and \\hcop\\ results also tend to agree to within an order of magnitude. \\item The overall blue excess of the sample of bright rimmed clouds is slightly less than that of the class~0 and class~I sources observed by \\citet{mmtwbg} and \\citet{gemm}, though the small number of bright rimmed clouds we observed does not make this difference statistically significant. A larger survey of bright rimmed clouds is required to determine if this is a significant finding. We do however make a case for the fact that the heating of the collapsing cloud by the adjacent HII region could dampen the infall signature, lowering the blue excess of bright rimmed clouds. \\item We observed outflows around 5 of the 7 bright rimmed clouds, including new detections of outflows around SFO~13 and SFO~25. These outflows appear to have similar properties to other outflows detected in millimeter and submillimeter emission. \\end{enumerate}  We do not see direct evidence of triggering in these sources. We can not determine if star formation was induced in these clouds or if we are seeing the collapse of pre-existing clumps. We do know that the environment has a profound effect on these regions. Although we have found similarities and differences between Bok globules and bright rimmed clouds, a detailed understanding of the effects of an ionization front on star formation can only be achieved by theoretically modeling this process, and then comparing that model to observations. In a forthcoming paper we will compare these observations with models of shock driven collapse derived by VC. In addition to bright rimmed clouds, these models could explain the effect of outflows and other environmental effects on star formation. Development of the techniques and models which will help us understand star formation in complex environments is in progress.  We gratefully acknowledge the staff of the HHT for their excellent support. In particular we wish to thank Harold Butner for his assistance with setting up our OTF observations at the HHT and Mark Heyer for his many helpful comments. C. H. De Vries and G. Narayanan are supported by the FCRAO under the National Science Foundation grant AST 01-00793."
    },
    "0206/astro-ph0206379_arXiv.txt": {
        "abstract": "We present a large sample of H- and K-band spectra of 32 optically line-luminous central cluster galaxies. We find significant rovibrational \\H2 emission in 23 of these galaxies as well as H recombination and/or [FeII] emission in another 5. This represents a fourfold increase in the number of molecular line detections known. A number of the detections are of extended emission (5--20~kpc). In several objects we find significant [SiVI] emission that appears to correlate with the strength of high ionization lines in the optical (e.g. [OIII]). This comprehensive sample builds on previous work and confirms that warm (1000--2500~K) molecular hydrogen is present wherever there is ionized material in the cores of cooling flows and in most cases it also coincides with CO molecular line emission. ",
        "introduction": "The cooling time of hot X-ray-emitting gas in the central regions of massive, relaxed clusters of galaxies can be substantially less than the Hubble time.  The gas in these regions can thus cool and recombine, initiating a cooling flow (Fabian \\& Nulsen 1977; Cowie \\& Binney 1977).  The ultimate fate of this cooling gas has been the subject of an extensive and strongly contested debate (see Fabian 1994). The cold gas has not been detected in molecular form and so is inferred to reside in a phase with $T_{\\rm gas} << 100$\\,{\\sc k}. Calculations of the gas properties are consistent with current observed limits (Ferland, Fabian \\& Johnstone 1994).  Until recently, the only cooling flow known to contain molecular gas was that around NGC~1275 in Perseus (see Bridges \\& Irwin 1998), although the interpretation of this source is complicated by the strongly varying nuclear component.  Moreover, the presence of the molecular gas may be related to the apparently on-going merger in this system, which has been the subject of a long-running debate (Van den Bergh 1977; Hu et al.\\ 1983; Pedlar et al.\\ 1990; Holtzman et al.\\ 1992; Norgaard-Nielsen et al.\\ 1993).  Another observational window on molecular hydrogen lies in the near-infrared where a number of strong rovibrational lines fall in the K-band. These lines have been detected in 8 different cooling flow central cluster galaxies (Jaffe \\& Bremer 1997; Falcke et al.\\ 1998; Krabbe et al. 2000; Wilman et al.\\ 2000; and Jaffe, Bremer \\& van der Werf 2001). The detection of molecular gas at 1000-2500~K implies excitation is required and that only a very small amount of `hot' molecular gas is present. So there is evidence for some molecular gas in the cores of cooling flows but its relation to the deposited gas predicted from the X-ray observations is far from clear.   This situation has been radically altered with the detection of CO line emission in 16 central cluster galaxies by Edge (2001). The molecular gas masses of 10$^{9-11.5}$ M$_\\odot$ imply that around 2--10\\% of the total deposited gas predicted from previous X-ray observations can be accounted for directly. However, with the dramatically reduced mass deposition rates found in {\\it Chandra} and {\\it XMM-Newton} observations (Allen, Ettori \\& Fabian 2001, Schmidt, Allen \\& Fabian 2001, Peterson et al. 2001), the molecular gas masses derived from CO observations are within a factor of three of that now expected. The mass of molecular gas appears to correlate best with the optical emission line luminosity indicating that the molecular gas observed is warmed from a  lower temperature. The rapid progress in this field has been possible due to the selection of massive cooling flows from the ROSAT All-Sky Survey (see  Crawford et al. 1999), few of which were known before 1990. To provide a direct comparison of the `hot' (1000--2500~K) and `cool' (30~K) molecular gas components requires near-infrared spectra of the systems that have been searched for CO.  In this paper we present the spectra and analysis for a systematic search for near-infrared hydrogen molecular lines in a complete sample of strong optical emission line emitting central galaxies drawn from Crawford et al. (1999). An ${H\\alpha}$ flux limit of $>3\\times10^{-15}$~erg~cm$^{-2}$s$^{-1}$ was adopted to ensure moderately bright lines and a redshift range of 0.03  to 0.3 to cover at least one \\H2 line in the K-band. This selection produces a sample of 18 objects reachable with UKIRT (one, A2146, is above the declination limit of 60$^\\circ$). Of this sample, only one object, A478, has not been observed in this study but it has a published spectrum (Jaffe, Bremer \\& van def Werf 2001). We also include seven weaker line emitting or lower redshift objects from Crawford et al. (1999) to fully sample any potential range in recombination lines to molecular line ratios, seven non-BCS central cluster galaxies with strong optical emission lines drawn from the literature (e.g. NGC~1275, A2597 and PKS~0745$-$191), and one control object with no lines (A2029) giving a total sample of 32 spectra. We also extend our spectral coverage to [FeII] in the majority of our targets making this the first comprehensive study of this line in central cluster galaxies.  Throughout we assume $\\Omega_0=1$ and $H_0 = 50$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "The study presented in this paper confirms that there is a remarkably good correspondance between the warm and cool molecular gas and  ionized hydrogen in intensity for central galaxies in cooling flows. These components appear to be common to this class of objects and indicate that substantial masses of irradiated molecular gas lie in the centres of clusters. What is required now is to establish the spatial correspondance between the observed components to these systems through narrow-band imaging and intergral field spectroscopy.  In the companion paper (Wilman et al. 2002) we present an analysis of the likely excitation mechanisms for the molecular lines observed.  The ultimate question of whether the total mass of cold molecular clouds deposited by the cooling flow is comparable to that predicted from X-ray observations cannot be addressed until observations of the far-infrared lines that should be emitted by the coolest clouds as they themselves cool. These lines are within reach of {\\it SOFIA, SIRTF} and {\\it FIRST} so this issue may be resolved in the next 2--4 years."
    },
    "0206/astro-ph0206186_arXiv.txt": {
        "abstract": "\\label{sec:abstract}  We develop the white dwarf luminosity function (LF) of the old open cluster NGC\\,188 in order to determine a lower limit to the age of the cluster by using the faint end of the cooling sequence.  To produce an extensive sequence of the cooling white dwarfs we imaged four contiguous HST-WFPC2 fields in the center of the cluster in the F555W and F814W filters.  After imposing selection criteria on the detected objects we found a white dwarf cooling sequence (down to V $\\simeq$ 26.5) including 28 candidate white dwarfs in the cluster.  The exposures are not deep enough to reach the end of this sequence, but the results of our analysis allow us to establish a lower limit to the age of the cluster independently of the isochrone fit to the cluster turnoff.  The most ancient white dwarfs found are $\\simeq$ 4 Gyr old, an age that is set solely by the photometric limit of our data.  Classical methods provide an estimate of $\\simeq$ 7 Gyr (Sarajedini et al., 1999). ",
        "introduction": "Stellar evolution theory predicts that all single stars having a Main Sequence (MS) mass lower than $\\simeq$ 8 M$_{\\rm \\odot}$ end their lives as white dwarfs (WDs).  In a star cluster, the white dwarf population carries key information for addressing a number of astrophysical questions. In particular, since the white dwarfs cool at a predictable rate, their LF can be used to constrain the cluster age. This age estimate can provide an independent test of the more traditional age determinations based on models and observations of the main-sequence turnoff (MSTO).  The white dwarf cooling sequence in the color-magnitude diagram, in combination with models can be used to estimate the initial mass -- final mass relationship between the stellar progenitors and the white dwarfs over the range of progenitor mass from near 8M$_{\\rm \\odot}$ to the current-day turnoff mass of a cluster. This allows a determination of the amount of mass lost from stars during their evolution through stellar winds and planetary nebula ejection.  Ideally, to employ white dwarfs as age estimators, we need to observe a large enough sample in any cluster to define the WDLF sufficiently well to pinpoint the luminosity of the turndown with a precision equivalent to a Gyr of cooling time. With smaller samples, it is still possible to make some progress on the age questions. For example, even a single good candidate WD in a cluster at magnitudes fainter than the faintest WDs predicted from the MSTO age and cooling theory can bring into question the vailidity of the WD {\\it or} MSTO-based ages.  Clearly the most interesting case for testing MSTO-based ages is for the globular clusters. However, because of the combination of large age and large distance for even the nearest of the Galactic globulars, it is not trivial to reach the end of the cooling sequence (e.g. Richer et al., 1995, 1997, 2002; Hansen et al. 2002). Old open clusters provide an interesting alternative for WD-cooling age studies. There are several populous clusters which are both nearer and younger than the nearest globulars and for which the expected WDLF turndown is 3 magnitudes or more brighter. For open clusters older than $\\sim$ 4 Gyr, the main-sequence turnoff stars are very similar in structure to globular cluster turnoff stars, most importantly in that they have radiative cores, and comparison of nuclear and cooling ages for these old open clusters are very relevant the issue of globular cluster MSTO age tests. A significant difficulty with open clusters is that even the most populous are sparse compared to most globular clusters. This leads to the problem that the lowest luminosity WDs are typically dramatically outnumbered by faint blue galaxies in the cluster fields. Attempts to statistically derive a WDLF by subtracting counts from nearby control fields are compromised by small errors in control-field counts and cosmic dispersion in faint blue galaxy surface density. Very accurate star-galaxy separation is therefore crucial for open cluster WD studies. Given the compact size of many faint blue galaxies, Hubble Space Telescope (HST) imaging gives a tremendous advantage. Ultimately another powerful technique for isolating cluster-member WDs from the compact blue galaxy and field WD populations may be measuring proper motions.   NGC\\,188 was considered for a long time to be the oldest observable open cluster in the Galaxy. Almost 40 years ago Sandage gave an estimate for the age of this cluster of 10 Gyr (Sandage, 1962; see also \\cite{eggen}).  Recent papers revise the early age estimates downward, and more modern values give an age near 7 Gyr (see Sarajedini et al., 1999). Other open clusters have now been found to be older than NGC\\,188 (NGC\\,6791 and Berkeley 17 with ages of 8 and 12 Gyr respectively, see \\cite{phelps}) however, NGC\\,188 is still one of the most studied since it is both nearer and less obscured than some of the other old open clusters.  Sects.  \\ref{sec:obs} and \\ref{sec:analysis} present the data, the reduction procedures and the color magnitude diagrams. The resulting LFs together with final results are presented and discussed in  Sects.  \\ref{sec:discussion} and \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Although we have identified a clear population of objects in NGC\\,188 with the photometric properties expected of cluster WDs, it is apparent from Fig. \\ref{fig:diff} that we have not reached the turnover in the WDLF. The comparison with the LF models suggests that we can place a lower limit of 5 Gyr on the age of the cluster.  These observations could be improved significantly with a modest effort, with HST and the ACS. A fainter photometric limit together with an expanded field coverage would allow a large sample of WDs to be identified in the NGC\\,188 field. NGC\\,188 stars have a proper motion of 2.3 milli-arcseconds with respect to the field and with a followup program, a subsample of bona-fide cluster WDs could be identified on the basis of their proper motions."
    },
    "0206/astro-ph0206465_arXiv.txt": {
        "abstract": "We report polarization measurements and photometry for the optical afterglow of the gamma-ray burst GRB 020405.  We measured a highly significant 9.9\\% polarization (in $V$ band) $1.3\\;$days after the burst and argue that it is intrinsic to the GRB.  The light curve decay is well fitted by a $t^{-1.72}$ power-law; we do not see any evidence for a break between 1.24 and $4.3\\;$days after the burst.  We discuss these measurements in the light of several models of GRB afterglows. ",
        "introduction": "\\label{sect_intro}   It is now commonly accepted that the radiation from gamma-ray bursts (GRB) afterglows comes from synchrotron radiation (e.g. M\\'esz\\'aros \\& Rees 1997). Since synchrotron emission is strongly polarized (e.g. Rybicki \\& Lightman 1979), one expects to see some level of polarization in GRB afterglows.  Several models make predictions regarding the amount of polarization one might see in a GRB afterglow (e.g. Gruzinov \\& Waxman 1999; Gruzinov 1999; Ghisellini \\& Lazzati 1999; Sari 1999). If the magnetic field is globally random but with a large number of patches ($\\approx 100$) with a coherent magnetic field, the polarization may be high, particularly in early times, depending on the number of patches (Gruzinov \\& Waxman 1999; Gruzinov 1999). Another alternative is that the field is globally symmetric (with no net polarization) but viewing the jet off-axis will break the symmetry and a significant polarization may result (Ghisellini \\& Lazzati 1999; Sari 1999). Microlensing could also explain a temporarily high polarization (Loeb \\& Perna 1998; Garnavich, Loeb \\& Stanek 2000).  On the observational side, polarization has been measured in several afterglows. \\citet{cov99} first detected 1.7\\% polarization for GRB 990510 about $0.77\\;$days after the burst, and \\citet{wij99} reported about the same amount of polarization for this burst $0.86\\;$days after the burst; the polarization angle being the same within the error bars. Several measurements for GRB 990712 indicated a polarization variability \\citep{rol00}, from $2.9\\% \\pm 0.4\\%$ to $1.2\\% \\pm 0.4\\%$ within $0.25\\;$days while the angle did not appear to change.  GRB 020405 was seen by several spacecrafts on 2002 April 05.028773 (UT) and IPN triangulation gave a tentative position of the burst \\citep{hurley}.  Prompt optical observations \\citep{pricea} revealed a new optical source at the position $\\alpha_{2000} = 13^h 58^m 03{\\fs} 1, \\delta_{2000} = -31{\\degr} 22\\arcmin 22\\arcsec$, which turned out to be quickly fading (Hjorth et al. 2002; Price et al 2002b; Covino et al 2002a). A redshift of $z=0.695$ was obtained with emission lines of the likely host galaxy \\citep{masetti}, later refined to $z = 0.6898 \\pm 0.0005$ \\citep{pricec}.  Polarization observations were reported by Covino et al. (2002a, 2002b, 2002c). ",
        "conclusions": "\\label{sect_disc}  The procedure described in Section~\\ref{sect_pol} resulted in measuring the flux difference between the MiniCam observations of the OT taken in May 2002 and those taken in April. This allows us to obtain the total flux of the OT and also the flux differences as measured through the polarization filters.  These flux differences, after subtracting a slight linear trend consistent with the decay of the OT, are shown in Fig.~\\ref{fig_pol}. As two polarization sequences were obtained, for each polarization filter two independent values of flux difference were measured. There is a clear polarization signature present, with the two sequences giving very similar values. Including the calibration uncertainty, a sinusoid fit yields $A=9.89\\pm 1.3$\\% for the polarization amplitude and $\\theta=-0.1\\pm 3.8 \\degr$ for the polarization angle, with resulting $\\chi^2=2.3$ for six remaining degrees of freedom (compared to $\\chi^2=89.0$ if zero polarization is assumed).  Assuming a $1/A$ prior (polarization cannot be negative) yields a very similar amplitude of $A=9.78\\pm 1.3$\\%. Fits to each individual sequence yielded $A=10.7\\pm 1.4$\\% and $\\theta = -0.4\\pm 4.2 \\degr$ for the first sequence, $A=8.5\\pm 1.7$\\%, $\\theta = 0.6\\pm 6.5 \\degr$ for the second sequence.  An immediate question is: could this measurement be caused by an instrumental effect? This can be checked by constructing identical flux difference curves for the stars in the field, which can be done easily when using the image subtraction method.  Most of the stars in the field have flux differences consistent with zero or relatively small ($<3$\\,\\%) polarization, and no other star shows the same polarization behavior as the OT.  In few of the cases when formally larger ($>3$\\,\\%) polarization is obtained from the fit, the polarization fits are not statistically acceptable, unlike for the GRB. We consider the detection of large ($\\sim 10$\\%) polarization in the OT of the GRB\\,020405 $1.31$ days after the burst to be very secure\\footnote{To allow the astronomical community an independent verification of this measurement, we have placed all relevant MMT MiniCam polarization data on {\\tt anonymous ftp} at {\\tt ftp://cfa-ftp.harvard.edu/pub/dbersier/GRB020405}.}.  Another question to answer is what is the chance that we would measure this level of polarization from random photometric errors alone? We created random flux measurements ($F_0,\\ F_{45},\\ F_{90},\\ F_{135}$) drawn from a a Gaussian distribution with average 1.0 and one-sigma width 0.02 (our error in magnitude). For each set of measurement we calculate the polarization with $$ P = 2\\times \\left[ \\frac{(F_0 - F_{90})^2 + (F_{45} - F_{135})^2}{(F_0+F_{45}+F_{90}+F_{135})^2} \\right]^{1/2} $$  This has been repeated ten million times. The first result is that for a photometric error of 0.03 mag, the most probable polarization is $2.0$\\%. The probability of having a polarization larger than 8\\% from photometric fluctuations alone is 1 in $10^7$ (for $\\sigma = 0.02$ mag); it is 1 in 1000 for $\\sigma = 0.03$ mag. However we have two consecutive sequences of measurements and they are both consistent with $>8$\\% polarization. Assuming that the errors in each sequence were 0.03 mag, the combined probability is $10^{-6}$. This does not account for the fact that the angle is the same for each of our observed sequences. The inescapable conclusion is that our polarization measurement is not due to chance photometric errors.  We calibrate the polarization to the surrounding star field so any Galactic effect would be due to material more distant than the surrounding stars.  Polarization could arise from the Galactic interstellar medium; however, the foreground reddening is so small [$E(B-V) = 0.05$] that it seems impossible that such a small amount of dust would produce such a large polarization. Dust in the immediate surroundings of the GRB could also induce some polarization. However, after correcting for the small Galactic reddening, the broad-band $BVRI$ spectrum is already very well described by a power law (Fig.~\\ref{fig_spectrum}), which is expected for a GRB. For interstellar polarization one has $P \\leq 9 E(B-V)$ \\citep{smf75}, which means that a color excess $E(B-V) \\sim 1.0$ would be needed to produce the observed amount of polarization. This would produce a intrinsic broad-band spectrum with a heretical slope of $+1.86\\pm 0.34$, and a very poor fit in any case. We thus consider that dust, whether in the host galaxy or in our Galaxy, is not the reason for the observed level of polarization."
    },
    "0206/astro-ph0206303_arXiv.txt": {
        "abstract": "{ Recently,   Norton   et  al.   (\\cite{Norton}),    on   the  basis   of multiwavelength photometry of \\wga\\,, argued that the\\ --1 day alias of the strongest peak in the power spectrum is the true orbital period of the system, casting doubts on the period estimated by Zharikov et al. (\\cite{Zharikov}).  We re-analyzed this system using our photometric and spectroscopic data  along with  the data kindly  provided by Andy Norton  and confirm our  previous finding.  After refining our analysis we find that the true orbital period of this binary system is 4\\fhour35. ",
        "introduction": "Israel et  al.  (\\cite{Israel1}) discovered that \\wga\\,  was a pulsating X-ray source. Strong modulations of  this source in X-rays were  obtained from the ROSAT PSPC ($721\\pm14$ sec)  and a more accurate period of $734\\pm1$ sec from ASCA was presented  by  Israel  et  al.   (\\cite{Israel1}) and  Israel  et  al. (\\cite{Israel2}).  Photometric observations of the optical counterpart of \\object{\\wga} exhibited  strong optical variations, compatible with the X-ray  (within 12 min) period (Uslenghi  et al.  \\cite{Uslenghi}). This modulation was interpreted as an evidence of the spin period of the WD in a close binary system.  Uslenghi et al. (\\cite{Uslenghi}) detected a circular  polarization from  the source  in the  R and  I  bands, with evidence  for a possible  modulation of  the polarization  at  twice the previously    observed   pulsation    period. \\wga was announced as  a new Intermediated  Polar (IP) by Negueruela  et  al.  (\\cite{Negueruela}) from  spectral  observations. Zharikov    et   al. (\\cite{Zharikov})  obtained   time  resolved spectroscopy and R-band   photometry from  which they deduced  an orbital period  of 4\\fhour36 and confirmed the pulsation period of 733\\,sec.  Later on, Norton et al., (\\cite{Norton})  obtained  UBVRI photometry  and  reported  that the  orbital period  was $5.387\\pm0.006$ hours, corresponding to the\\  --1  day alias of the  period found by Zharikov  et  al.  (\\cite{Zharikov}). They  had  some  ambiguity  in determining  which  of  the  daily  cycle  aliases  of  low (orbital) frequency and  intermediate (beat) frequency   to  pick  up,  because selecting the strongest  peak in low frequencies was  forcing the beat period  into  a\\ --2  day  alias  of  the intermediate  frequency  peak. Through detection  of the beat  frequency, Norton et  al. (\\cite{Norton}) also confirmed that the rotational  period of the white dwarf is twice the  pulse  period,  and  they  confirmed  the  presence  of the circular polarization in the source by detecting oppositely signed polarization in each of the B and R bands.  In  this  letter,  we  re-analyze  our spectral  and  photometric  data together with photometric data from Norton et al. (\\cite{Norton}) confirm and refine our previous period estimate of 4\\fhour35. ",
        "conclusions": "Norton  et  al. (\\cite{Norton})  chose  the $\\Omega_{\\mathrm N}  =  4.455\\pm0.005$\\ $\\mathrm d^{-1}$,  or $P_{\\mathrm N}=5.387\\pm  0.006\\ \\mathrm h$,  as  the orbital  period of  the system from the analysis of the  power spectrum  peak strength combination. They noted  that the power  spectrum is dominated by  three sets of signals at $\\sim 5.5\\ \\mathrm d^{-1}$,  $55.5\\ \\mathrm d^{-1}$ and $117.8\\ \\mathrm d^{-1}$ but the strongest peaks in each of the three sets are not harmonically related to  each  other.   The   solution  $\\Omega_{\\mathrm N}$  was  selected  as the  more probable. They assume that more extreme aliases combinations are  unlikely,  since  the  power at  these  alias are  low, although  such combination  are not excluded. In our opinion the strength of peaks of power spectra are highly dependent on the quality of the data and sampling. The photometric data of Norton et al. (\\cite{Norton})  is certainly undersampled for such far-reaching conclusions. On the other hand, the spectroscopic observations presented here unambiguously identify the orbital period of the system.    Adding  the  data kindly  provided  by  authors  of Norton  et  al. (\\cite{Norton}) to our measurements, we were able to improve slightly the period estimate.   The new  value for the  period of  the Intermediate Polar \\wga\\, now stands at $4\\fhour35\\pm0\\fhour01$,   similar to our recently reported  value (Zharikov et al.  \\cite{Zharikov}).  We note that this analysis does not  change our previous estimates of the system parameters,  but shifts  the photometric  minimum in  the  light curve exactly to  the redefined  epoch $T_{\\mathrm 0} =  2451762.9527\\pm0.0001$, which corresponds to the\\ \\ $\\pm$ zero crossing of the \\ion{H$_\\beta$}{} radial velocity  curve, i. e.   to the  moment when  the secondary  is located between the observer and  the WD.  The final phase-folded  light curves in the R band,  AFD, and radial velocity curves  in \\ion{He}{ii} 4686 and \\ion{H$_\\beta$}{} are presented in Figure \\ref{fig5}. The difference of amplitudes and phases of the \\ion{H$_\\beta$}{} and  \\ion{He}{ii} lines were discussed in our previous paper.    \\begin{figure}[t] \\includegraphics[width=8.5cm,clip=]{ED112_Fig5.eps} \\caption{The  radial  velocity  curves   of \\ion{H$_\\beta$}{}  and \\ion{He}{ii} 4686,  folded with  the spectroscopic orbital   period  of $4\\fhour35$,   are   presented in the  middle   panel. The combined   $R_\\mathrm{c}$  light   curve  of   \\object{\\wga}  is presented in the lower panel.  The  data of Norton et al.  (\\cite{Norton}) is marked with  open  circles. Full circles  are  from  Zharikov et al. (\\cite{Zharikov}). The AFD (all filter data) folded in the same manner is shown in the top panel.} \\label{fig5} \\end{figure}"
    },
    "0206/astro-ph0206135_arXiv.txt": {
        "abstract": "{ We present results from a survey of molecular hydrogen emission from a sample of Starburst and Seyfert galaxies carried out with the Infrared Space Observatory (ISO). Pure rotational H$_{2}$ emission has been detected in a number of extragalactic objects and a variety of environments. A number of transitions from S(7) to S(0) are detected in both Starbursts and Seyferts. Using excitation diagrams we derive temperatures and masses of the ``warm'' molecular hydrogen. We find that the temperature of the ``warm'' gas is similar in Starbursts and Seyferts (those Seyferts for which we have firm detections of the S(0) line) with a value of around T$\\sim$150 K. This ``warm'' gas accounts for as much as 10\\% of the total galactic mass (as probed by CO molecular observations) in Starbursts. The fraction of ``warm'' gas is overall higher in Seyferts, ranging between 2--35\\%. We then investigate the origin of the warm H$_{2}$ emission. Comparison with published theoretical models and Galactic templates implies that although emission from photodissociation regions (PDR) alone could explain the emission from Starbursts and Seyferts, most likely a combination of PDR, shock emission and gas heated by X-rays (mostly for the Seyferts) is responsible for H$_{2}$ excitation in extragalactic environments. Finally, we find that although PAH and H$_{2}$ line emission correlate well in Starbursts and the large scale emission in AGN, H$_{2}$ emission is much stronger compared to PAH emission in cases where a ``pure'' AGN dominates the energy output. ",
        "introduction": "Using the Short Wavelength Spectrometer (SWS; De Graauw et al. 1996) on board the Infrared Space Observatory (ISO; Kessler et al. 1996), we have performed a survey of molecular hydrogen emission from active galaxies displaying a wide range in nuclear activity including pure bona-fide Starbursts, Seyfert 2s (some of them with starburst components) and pure Seyfert 1 galaxies. Prior to the ISO mission, extragalactic H$_{2}$ emission had only been detected in the ro-vibrational lines around 2.1 $\\mu$m. Indeed, H$_{2}$ ro-vibrational emission has been detected in galactic sources (e.g. Usuda et al., 1996), Starburst galaxies (e.g. Joseph, Wright and Wade 1984), Seyferts (e.g. Moorwood and Oliva 1988, Fischer et al. 1987 ) and bright spirals (e.g. Puxley, Hawarden and Mountain 1988).   The ro-vibrational lines typically trace H$_{2}$ gas of masses around $\\sim$10$^{4-5}$ M$_{\\odot}$ and temperatures $\\sim$2000 K. This gas can be excited either by collisions (thermal) or by absorption of ultraviolet (UV) photons in the Lyman and Werner electronic bands (912-1108 \\.A), followed by a de-excitation cascade to the ground state (fluorescence). However, gas at these temperatures is a very small fraction (as small as 10$^{-6}$) of the total amount of H$_{2}$ gas (e.g. Van der Werf et al. 1993). Since the ro-vibrational lines tend to get faint at lower temperatures, most of our knowledge about the H$_{2}$ content of galaxies comes from CO observations assuming a CO/H$_{2}$ conversion factor derived from galactic molecular cloud observations. ISO gave the unique opportunity to observe intermediate temperature gas, ie ``warm'' H$_{2}$, directly in pure rotational lines. Since transitions with $\\Delta$J= $\\pm$ 1 are strictly forbidden for the H$_{2}$ molecule, the rotational ladder consists only of an ortho (J odd) and a para (J even) series of quadrupole transitions.   ISO and in particular SWS, offered the unique opportunity to detect pure rotational H$_{2}$ emission in a number of galactic and  extragalactic sources, thus studying the amount of moderately warm gas in these sources. The spectral range of SWS provides full coverage of a number of transitions (for most galaxies we have observed from the S(0) to S(7) transitions) while its spectral resolution is well matched to the typical velocity dispersions of galaxies.  Among the first detections of pure rotational H$_{2}$ emission in galaxies were the detections in NGC 3256 (Rigopoulou et al. 1996), NGC 6946 (Valentijn et al. 1996), NGC 891 (Valentijn \\& Van der Werf 1999). However, no study of pure rotational H$_{2}$ emission for a large number of galaxies has so far appeared.  Here, we present an inventory of H$_{2}$ emission lines from a number of Starburst and AGN. Our survey includes 12 and 9 Starburst\\footnote {The two components of NGC 3690 have been observed separately NGC 3690A and NGC 3690BC. In the analysis we treat them as two separate systems.} and Seyfert galaxies, respectively. Temperatures of the warm molecular gas are deduced from excitation diagrams whereas the masses of the warm molecular gas are compared to the total gas content of the galaxies as estimated from molecular CO observations. The H$_{2}$ excitation mechanism is investigated next. The observations are compared to the predictions of published models both for PDR, shocked emission and X-ray irradiated gas, as well as to Galactic templates. Finally we examine possible correlations between PAH and H$_{2}$ emission. ",
        "conclusions": "We have presented pure rotational H$_{2}$ emission lines observed with the ISO satellite from a sample of Starburst and Seyfert galaxies. We have compared the emission properties of the two samples. The results are summarized as follows:\\\\  The lowest S(0) transition has been detected in 5 out of the 10 Starbursts and 4 out of the 9 Seyfert galaxies. The S(1)$/$S(0) line ratio is within the errors not that different in Starbursts and in Seyferts.  The temperature of the ``warm'' gas, as estimated using the S(0) detections was found to be similar (within errors) T$\\sim$150 K, in both Starburst and Seyfert galaxies. This in turn implies that the global properties of the gas in both environments are the same. We caution though that due to the large ISO beams dilution of a pure AGN-related effect is important.\\\\  The ``warm'' gas mass constitutes up to 10\\% of the total molecular gas content (as traced by CO molecular observations) in Starbursts. However, the fraction of ``warm'' gas in Seyferts is considerably higher reaching up to 35\\% (although with a large scatter in values). We propose that, extra ``heating'' of the gas is provided by energetic hard X-ray photons originating from the central AGN.   We have compared the observed strength of the molecular H$_{2}$ lines with theoretical model predictions as well as those of local Galactic templates.  Such comparisons reveal that a combination of various PDR clouds explain  reasonably the line ratios observed in Starbursts. Although such a combination of PDR models can also match the observed line ratios in Seyferts it is likely that slow velocity shocks and some heating from the central X-rays are also present.  Finally, we have examined the existence of a link between PAH and molecular line emission. We find that Starbursts and Seyferts with a strong starburst-component follow a very similar correlation. On the other hand the ratio H$_{2}/$PAH is higher in pure ``AGN--dominated'' objects. It is likely that an extended circumnuclear component of ``warm'' gas (heated by the nuclear X-ray emission) is present in AGN in which enhanced H$_{2}$ emission originates. This ``warm'' gas component is either too far away for UV photons to reach or, acts as a shield to UV photons resulting in both cases in suppressed PAH emission."
    },
    "0206/astro-ph0206245_arXiv.txt": {
        "abstract": "s{ We discuss the present status and future prospects for cosmic shear observations and their cosmological constraints. We review the evidences supporting the cosmological origin of the measured signal, and discuss the possible problems coming from intrinsic alignment and the actual limitations of theoretical predictions. } ",
        "introduction": "The cosmic shear is a gravitational lensing effect which occurs everywhere in the universe, and allows astronomers to map the projected mass distribution on the sky from the solely observation of the distorted distant galaxies. The idea of mapping the matter using the gravitational deflection of ray-lights was born in 1937 when F. Zwicky \\cite{zwa,zwb} envisioned the possibility to use the distorted shape of distant galaxies to probe the matter content in nearby clusters of galaxies. His  idea was only discussed seriously 30 years after, when first detailed analytical work were produced \\cite{ks,gunn}, motivated by the progress made in geometric optic in curved spacetimes \\cite{s61}. But it is only in the early 90's that a robust link was established between an appealing theoretical idea and the observational possibilities \\cite{me91,b91,k92}. In the meantime, in 1983, a first attempt to measure cosmic shear failed \\cite{v83} mainly because of the poor image quality of data available at that time. The interest for cosmic shear raised again after the discovery of giant arcs in 1987, and the burst of theoretical papers started in the mid 90's \\cite{v96,b97,j97}, followed by many others \\cite{bs01}. In 2000  the first detections were reported almost simultaneously, and independently by four teams \\cite{b00,k00,vw00,w00}. Since then, several other measurements were done and significant improvements in the data analysis lead to refined measures and to the first robust cosmological constraints \\cite{maoli01,vw01,rhodes01,pen02,vw02,hoekstra01,ham02,hoekstra02,bacon02,ref02}. ",
        "conclusions": ""
    },
    "0206/astro-ph0206073_arXiv.txt": {
        "abstract": "A 0.2-12 keV spectrum obtained with the \\xmm\\ EPIC/pn instrument of \\grb, taken in the first 5~ksec of a 27 ksec observation, was found by Reeves \\etal\\ (2002; R02) to contain emission lines which were interpreted to be from \\magnesium, \\silicon, \\sulfur, \\argon, and \\calcium, at a lower-redshift ($z_{obs}=1.88$) than the host galaxy ($z_{host}=2.14$).  We examine the spectrum independently, and find that the claimed lines would not be discovered in a blind search. Specifically, Monte Carlo simulations show that the significance of reported features, individually, are such that they would be observed in 10\\% of featureless spectra with the same signal-to-noise. Imposing a model in which the two brightest lines would be \\silicon\\ and \\sulfur\\ K$\\alpha$ emission velocity shifted to between $z$=1.88--2.40, such features would be found in between $\\sim$1.3-1.7\\% of observed featureless spectra (that is, with 98.3-98.7\\% confidence).  When we account for the number of trials implicit in a search of five energy spectra (as were examined by \\rr), and permit a wider $z$-phase space search ($z=2.14$\\ppm1.0), the detection confidence of the two line complex decreases to 77-82\\%.  We find the detection significances to be insufficient to justify the claim of detection and the model put forth to explain them.  $K\\alpha$ line complexes are also found at $z=1.2$ and $z=2.75$ of significance equal to or greater than that at $z=1.88$.  Thus, if one adopts the $z=1.88$ complex as significant, one must also adopt the other two complexes to be significant. The interpretation of these data in the context of the model proposed by \\rr\\ is therefore degenerate, and cannot be resolved by these data alone. Our conclusions are in conflict with those of \\rr, because our statistical significances account for the multiple trials required -- but not accounted for by \\rr\\ -- in a blind search for emission features across a range of energies.  In addition, we describe a practical challenge to the reliability of Monte Carlo \\dchisqr\\ tests, as employed by R02. ",
        "introduction": "It was recently reported that the X-ray afterglow of gamma-ray burst \\grb, as observed with \\xmm\\ EPIC/pn, contained spectral emission lines \\cite[hereafter, \\rr]{reeves02} -- the first report of multiple X-ray emission lines from a {\\small GRB}.  These lines, at 0.45, 0.70, 0.89, 1.21, and 1.44 keV were interpreted to be from He-like \\magnesium\\ (rest energy 1.35 keV) and H-like \\silicon (2.0 keV), \\sulfur\\ (2.62 keV), \\argon\\ (3.32 keV) and \\calcium\\ (4.10 keV), redshifted to $z$=1.88.  The difference between this and the known redshift of the host galaxy $z_{host}=2.14$ was modeled as due to supernova ejecta traveling at $v=25800$\\ppm1200 km \\perval{s}{-1}, which had originated during a supernova 4 days prior to when the {\\small GRB} jet illuminated it, producing the afterglow (a more detailed analysis by the same authors was completed after this paper was in its initial form; \\citenp[\\rrb\\ hereafter]{reeves02b}  The statistical significance of the individual lines was not reported in \\rr; it was stated that joint analysis of the lines taken together produced an improvement in the \\chisqr\\ value which, by an F-test, yielded a significance level of 99.7\\%.  In addition, it was found that Monte Carlo (MC) simulations were unable to produce the the same improvement in \\chisqr\\ found between the best-fit power-law model and the best-fit five emission-line model more than 0.02\\% of the time. Specifically, it was found that the best-fit \\chisqr\\ value for a power-law model was improved by fitting to a model of a MEKAL plasma with emission lines at rest energies corresponding to unresolved \\magnesium, \\silicon, \\sulfur, \\argon, and \\calcium\\ redshifted to $z=1.88$ in only 0.02\\% of the simulated spectra (a 99.98\\% confidence detection).  The implications of the model discussed by \\rr\\ -- a delay between a supernova and a {\\small GRB} on a timescale of days, the formation of a thin shell of supernova ejecta, an apparent under-abundance of Fe relative to the detected nuclei -- provide severe constraints on gamma-ray burst emission models.  In addition, as demonstrated by \\rr, the future detection of multiple emission lines can provide extremely strong constraints on the production mechanisms, due to the inherent required outflow velocity and emission timescales which can be derived from them, not to mention the implied association with supernovae. Similar spectra observed with greater S/N in the future would greatly aid in unravelling the emission mechanisms and geometry of gamma-ray bursts.  Therefore, it is of wide theoretical (e.g. \\citenp{lazzati02,kumar02}) and observational interest to further interpret the observed X-ray spectrum of this \\grb, in hopes of determining what more could be learned from future, more precise observations.  In Sec.~\\ref{sec:anal}, we describe the observation, and perform a basic spectral analysis using continuum models.  In Sec.~\\ref{sec:mc}, we compare Monte Carlo (MC) realizations of acceptable continuum models with the {\\small GRB} spectrum, and find that the reported features would be produced in $\\sim$10\\% of the continuum model spectra, due only to Poisson noise.  In Sec.~\\ref{sec:ka}, we adopt the model that the two apparently most significant lines are K$\\alpha$ lines of \\silicon\\ and \\sulfur; we perform a blind search for features of the same significance in MC realizations of continuum spectra, and find that they would be reported from $\\sim$1.2-2.6\\% of such spectra, again, due only to Poisson noise.  We describe in Sec.~\\ref{sec:mcbad} a practical challenge to the reliability of the MC \\dchisqr\\ analysis produced by R02.  We conclude in Sec.~\\ref{sec:con} that the lines are not individually significant in the absence of an imposed model, and are only marginally significant when the adopted model is imposed.  {\\change\\ These conclusions conflict with those of \\rr. R02 derived the model (that is, the observed line energies, or redshift) from the data, and then applied statistics for detection as if the energies were known prior to examining the data (that is, single-trial statistics). This is not appropriate when the model line energies are derived directly from the X-ray data, and not from an a priori model -- one derived without examination of the X-ray data (for example: line energies of multiple features with redshifts of the host galaxy).  We adopt statistics appropriate to a blind-search for these features, across a range of energies or redshifts (multi-trial statistics).  This accounts for the diminished significance we find for the features. } We further discuss the reasons for this conflict and conclude in \\S~\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  We have attempted to confirm the observational statistical significance of emission lines in the X-ray afterglow of \\grb.  In a blind-search for individual emission lines between 0.4 and 1.5 keV, features of significance equal to those observed will be found in one in ten featureless spectra.  Thus, the reported features can be said to be detected with 90\\% confidence in a model-independent way.  Also, a blind-search for the two-line complex (\\silicon\\ and \\sulfur) at any redshift between the reported value ($z=1.88$) and a blueshift of equal magnitude from the host galaxy ($z=2.40$) would find such features with equal significance to that observed in 1 of 60 featureless spectra (1.3-1.7\\% of the time, depending on the intrinsic spectrum).  Thus, the features as reported can be said to be detected with 98.3-98.7\\% confidence, in a model-dependent interpretation, where we search for two features due to \\silicon\\ and \\sulfur\\ K$\\alpha$ redshifted to some value of $z$ in the range $z=$2.14\\ppm0.26.  The difference between the present statistics and those of \\rr\\ are due to the different statistical arguments used to establish the existence of the emission lines.  While \\rr\\ relies on single-trial statistics, we find the model used by \\rr\\ (K$\\alpha$ lines, at a redshift different from that of the host galaxy) was derived directly from the data, which therefore requires a statistical analysis appropriate to a blind search.  By expanding the searched phase-space, and taking into account the multiple trials of a blind search, the confidence in the detection drops from the 99.98\\% of \\rr\\ to, the 98.7\\% (best case) we find here.  Moreover, \\rr\\ did not estimate the individual significances of the lines as we do here; thus we find that such ``lines'' would appear in between 15-78\\% of observed featureless spectra for a single-trial significance comparable to that of the reported \\silicon\\ or \\sulfur\\ lines.  { The analysis of these data has otherwise recently been called into question.  \\acite{lanl02} have shown that there is a background line associated with the EPIC/pn detector edge during the observation, which would have been included in the {\\small GRB} spectrum from the first 5 ksec, when the source was near the detector edge, but not afterwards, after the source had been moved away from the detector edge.  In our own analysis, we cannot confirm this result unless we adopt non-standard event selection criteria, which differ from the ones used by \\rr.  \\rr\\ removed events near the CCD detector edge ({\\tt FLAG==0}) and selected only single and double events ({\\tt PATTERN<=4}) (J. Reeves, priv. comm.).  These selections result in a smooth, featureless background spectrum with with no bright line-like feature near $\\sim 0.7 ~\\rm{keV}$ (as seen in Fig.~4f of \\citenp{lanl02}) as well as a reduction of the count rate by a factor of $\\ga 2$ in the range $E = 0.2 - 3 ~\\rm{keV}$ (see Fig.~\\ref{fig:masao}).  Therefore, we are not able  to confirm the applicability of \\cite{lanl02} to the analysis of \\rr. }   An alternative approach to the one we have taken is employed using XSPEC, in which one fits a featureless spectrum to the data, and then a spectrum which includes emission lines, to determine if the change in \\chisqrnu\\ is significant, as according to an F-test; this is the approach taken by \\rr.  However, this approach for the detection of emission or absorption lines is formally incorrect, and gives false statistical results \\cite{protassov02} particularly so when the true continuum is not well constrained, as in the present case.  We therefore prefer our approach of applying a matched energy response filter for line detection at arbitrary energies, and to compare this with application of the matched filter to MC realizations of featureless spectra.  It is a trivial statistical exercise to demonstrate that matched filtering maximizes the signal-to-noise ratio (and thus detectability) for detection of infinitely narrow emission lines.   In estimating the model-dependent confidence limit for the detection of the line complex (98.7\\%), we accounted only for searching the redshift phase space between $z=1.88$ and $z=2.40$, symmetric about the host galaxy redshift -- an extremely minimal requirement.  We did not account for the full redshift phase space searched by \\rr , as such was not given in that reference; if the redshift phase space searched by \\rr\\ covered $z=1.14-3.14$ ($z=2.14\\ppm1.0$), then the detection significance of the two strongest lines (\\silicon\\ and \\sulfur) together decreases from 98.3-98.7\\% to 95-96.2\\% confidence. Finally, we did not include in this confidence limit the number of trials implicit in searching five X-ray spectra for emission lines, which was performed by \\rr\\ for different time periods (0-5 ksec, 5-10 ksec, 10-15 ksec, 15-20 ksec, and 20-27 ksec).  If we presume the same search was made on all five spectra, as seems a reasonable {\\em a priori} search to perform, then the detection confidence for the \\silicon\\ and \\sulfur\\ lines together decreases to $0.95^5$--$0.962^5$=77-82\\%.  We regard 98.7\\% to be a conservative (in the sense of permitting a higher significance) upper-limit to the confidence of detecting the \\silicon\\ and \\sulfur\\ lines together, while a more accurate accounting of the number of trials and phase-space searched by \\rr\\ produces a 77-82\\% confidence limit.  We consider neither a 90\\% confidence detection in a model-independent interpretation, nor a 98.3-98.7\\% confidence detection in a model-dependent interpretation, to be sufficient to justify the detection claims and subsequent interpretation put forth by \\rr.  The 77-82\\% confidence limit, which accounts for the wide $z$-phase space and number of spectra examined by \\rr, is well below any comfortable detection confidence. If the $z$ phase space actually searched by \\rr\\ is larger, the number of implicit trials is greater, and our estimate of the confidence level for the detected line complex would decrease.   Moreover, if one concludes that the marginal detection of the 2 lines (Si \\& S) near $z=1.88$ is significant, then one must also conclude that the detection of all 5 lines near $z=2.75$ is equally significant.  In addition, if one concludes that the marginal detection of the 5 lines near $z=1.88$ is  significant, then one must also conclude that the detection of 2 lines (Si \\& S) near $z=1.2$ is equally significant.  Therefore, one cannot conclude simply that a complex of K$\\alpha$ line emission is detected near $z=1.88$; these data permit alternate interpretations of such complexes near $z=1.2$ and $z=2.75$.  As the statistical excesses are due to the same ``features'' in the observed spectrum, the interpretation of the statistical excess in the context of the model presented by \\rr\\ is degenerate and cannot be resolved with these data alone.  Prospects for confirmation of line features in GRBs are very good, considering that the X-ray spectral integration for \\grb\\ was begun 11 hours after the {\\small GRB} was initially detected, and required 1.4 hrs of integration to obtain.  Decreasing the reaction time would permit a longer integration, while the afterglow is brighter in the X-rays, and the marginal results found here may well be improved upon."
    },
    "0206/astro-ph0206150_arXiv.txt": {
        "abstract": "{ We present a study of the three ultraluminous infrared galaxies IRAS\\,14348-1447, IRAS\\,19254-7245, and IRAS\\,23128-5919, based on mid-infrared (MIR) spectro-imaging (5--18\\,$\\mu$m) observations performed with ISOCAM. We find that the MIR emission from each system, which consists of a pair of interacting late type galaxies, is principally confined to the nuclear regions with diameters of 1--2\\,kpc and can account for more than 95\\% of their IRAS 12\\,$\\mu$m flux. In each interacting system, the galaxy hosting an active galactic nucleus (AGN) dominates the total spectrum and shows stronger dust continuum (12--16\\,$\\mu$m) relative to the Unidentified Infrared Band (UIB) emission (6--9\\,$\\mu$m), suggestive of its enhanced radiation field. The MIR dominant galaxy also exhibits elevated 15\\,$\\mu$m/H$\\alpha$ and 15\\,$\\mu$m/K ratios which trace the high extinction due to the large quantities of molecular gas and dust present in its central regions.  Using only diagnostics based on our mid-infrared spectra, we can establish that the Seyfert galaxy IRAS\\,19254-7245 exhibits MIR spectral features of an AGN while the MIR spectrum of the Seyfert (or LINER) member of IRAS\\,23128-5919 is characteristic of dust emission principally heated by star forming regions. ",
        "introduction": "It is currently widely accepted that the majority of the most luminous galaxies (L$_{bol}>10^{11}$\\,L$_{\\sun}$) in the local universe (z $<0.3$) are luminous in the infrared, and include the ultraluminous infrared galaxies (ULIRGs, L$_{\\rm IR}>10^{12}$L$_{\\sun}$) which emit the bulk of their energy at infrared wavelengths \\citep[ and references therein]{Houck1984,Soifer1989,Sanders1996}. In those systems most of the infrared emission seems to originate from their dusty nuclear regions. Even though one of the principal heating mechanisms for the lowest luminosity ($\\lesssim 10^{11}$ L$_{\\sun}$) infrared galaxies is the stellar radiation field of young massive stars, it is still unclear if the star formation is also the dominant heating source for ULIRGs or whether one needs to invoke an active galactic nucleus (AGN) and its strong radiation field as the central engine responsible for the heating of the dust \\citep[see][]{Joseph99,Sanders99}. The presence of large quantities of molecular gas has long been detected in the central regions of most ULIRGs \\citep[e.g.][]{Sanders1985,Sanders1991} leading to high extinction of both their UV and optical radiation. As a result, since it appears that most galaxies do harbor a super-massive, though often quiescent, black hole \\citep{Richstone1998}, one would expect to find in their galactic nucleus observational evidence for a mixture of AGN \\citep{Sanders1988} and/or strong compact starburst regions \\citep{Condon1991} fueled by the high concentration of molecular gas \\citep{Bryant1999}. Observations in the mid-infrared (MIR), which are less affected by absorption than shorter wavelengths \\citep[A$_{15\\,\\mu m}$ $\\sim$\\,A$_{V}$/70,][]{Mathis1990}, thus provide a powerful probe of galactic central regions \\citep{Soifer2000,Soifer2001}.  As we discussed in \\citet{Laurent2000}, the integrated MIR emission in active galaxies is produced mainly by the interstellar dust which is heated directly by the ionization field from young stars or an AGN. This is in contrast to late type galaxies where the MIR (5--20$\\mu$m) energy budget is dominated by the reprocessed emission of star forming regions in their disk and accounts for $\\sim$\\,15\\% of their luminosity \\citep{Dale2001,Helou2001,Roussel2001}. However, the main difficulty in assessing the importance of the underlying physics in galactic nuclei, where the spatial resolution is typically poor, is in separating the contribution of star forming regions and the active nucleus from the integrated MIR emission. The development, application, and general utility of MIR diagnostics in nuclei of galaxies has already been demonstrated by \\citet{Roche1991} and more recently by \\citet{Genzel1998,Laurent2000}, as well as by \\citet{Dudley1999,Imanishi2000}. This was mainly accomplished with the advent of ISOCAM and SWS on board ISO, with high spatial and spectral resolution, as well as improved sensitivity in the 3 to $\\sim$40\\,$ \\mu$m wavelength range, thus allowing us to study the nature of the heating sources in ULIRGs. More specifically it has been shown by \\citet{Lutz1998,Laurent1999b,Laurent2000,Tran2001} that a nearby galaxy hosting a dominant AGN is clearly different in the MIR from a starburst or a late type spiral. The most striking difference is that the rather featureless MIR spectrum in AGN lacks the emission bands at 6.2, 7.7, 8.6, 11.3 and 12.7\\,$\\mu$m, which are seen in late type galaxies and are attributed to Polycyclic Aromatic Hydrocarbons (PAHs) -- also often called Unidentified Infrared Bands (UIBs). One may consider that this is simply due to the fact that its elevated MIR continuum of the AGN overwhelms any UIB feature emission \\citep{Pier1992,Barvainis1987}. It seems inevitable that as the AGN heats its dusty torus at T$\\sim$\\,1000\\,K and the dust grains approach sublimation temperatures, the more fragile molecules responsibly for the UIB emission could be partly destroyed by a photo-thermo-dissociation mechanism \\citep{Leger1989}. Obviously this picture is more complicated in distant galaxies since due to limited spatial resolution the contribution of the star forming regions surrounding an AGN would progressively enter into the beam and dilute any AGN MIR signature \\citep[see][]{Laurent1999b}. When sufficient spatial resolution is available to directly view the active nucleus, as is often the case in Seyfert 1 galaxies, the non-thermal emission from the AGN will dominate the spectrum. Consequently, the spectrum can then be fitted by a power law and has a ``bump'' in the 4--5$\\mu$m range. A 5--11\\,$\\mu$m study of a large sample of Seyfert galaxies with ISO by \\citet{Clavel2000} confirmed this picture, concluding that Seyfert 2 galaxies have weaker MIR continuum. However, a detailed analysis of the MIR spectra and images of the prototypical Seyfert~2 galaxy NGC\\,1068 by \\citet{LeFloch2001} showed that if sufficient spatial resolution is available and the AGN is extremely strong, even in the case of a Seyfert 2 one can isolate the emission of the central engine from the star forming regions which surround it. In that case the MIR spectrum of the Seyfert 2 would also be a power law with the addition of a weak PAH emission.  Despite this progress, several questions concerning the extent and spectral characteristics of the MIR emission in active nuclei, as well as the correlation between MIR and optical activity have not been fully examined. Could broad band MIR photometry be used to probe the physical characteristics of AGNs? In the present paper we try to address some of these issues by studying the MIR spectral energy distribution (SED) of three ultraluminous IRAS galaxies. Each IRAS source, the properties of which are presented in Table~\\ref{info}, consists of a merging pair of galaxies with different levels of nuclear activity. The targets were specifically selected as MIR bright and harboring an optically classified AGN. In section 2, we describe the observations and in section 3 we present the details of our study and analysis of the data for each system. A discussion followed by concluding remarks is presented in section 4. Throughout this paper we assume a Hubble constant H$_{0}$=75 km\\,s$^{-1}$\\,Mpc$^{-1}$ and q$_0$=1/2. ",
        "conclusions": "\\label{sec:discuss}  A wealth of observational data available has shown that ULIRGs have high concentrations of gas and dust in their nuclei, sufficient to account for most of their observed infrared luminosity \\citep[see][ for a review]{Sanders1996}. Whether the energy source of ULIRGs is a dust enshrouded AGN or a starburst still remains an open issue. However, recent indirect evidence is beginning to favour the existence of bright extremely red point-like sources in the nuclear regions of ULIRGs. More specifically near-infrared observations of luminous infrared galaxies have shown that their flux at 2.2\\,$\\mu$m is more concentrated towards the center than at 1.3\\,$\\mu$m \\citep{Carico1990b,Scoville2000}. Furthermore, recent high resolution MIR observations using Keck of a sample of ULIRGs reveal compact sub-arcsecond sources (with linear scales of $\\sim$\\,100--300~pc) which contain 30\\% to 100\\% of the observed MIR energy of these galaxies \\citep{Soifer2000}. This contrasts with the LIRGs ($10^{11}$\\,L$_{\\sun}$ $\\leq$ L$_{\\rm IR} \\leq 10^{12}$\\,L$_{\\sun}$), in which the infrared energy seems to be generated over somewhat larger scales \\citep[$\\sim$\\,100\\,pc\\---1\\,kpc,][]{Soifer2001} and sometimes can be found in extra-nuclear regions associated with the physical interaction of merging pairs of galaxies. Furthermore, there are galaxies such as VV\\,114 where it has even been found that a substantial fraction of the MIR flux originates from an extended component of hot dust emission spread over several kpc scales \\citep{Soifer2001,LeFloch2002}. ULIRGs are thus not simply a scaled-up version of LIRGs and require further dynamical compression of the molecular gas responsible for the IR luminosity within very compact regions. A plausible mechanism would be one where the shocks and tidal forces of the interaction first lead to star formation over galactic scales, leading to IR luminosities up to a few 10$^{11}$\\,L$_{\\sun}$. Subsequently, gravitational instabilities and the formation of a bar, strip the gas of its angular momentum, funneling large quantities towards the nuclear regions of galaxies, which can feed circumnuclear starbursts or AGNs and trigger the ultraluminous phase in the infrared \\citep{Combes2001}.  Even though the above scenario is appealing, given the high extinction in the nuclei of ULIRGs, the limited atmospheric transmission in the MIR windows, and the limited sensitivity of ground-based instruments, questions related to the direct probing of the nuclear activity such as ``does all MIR emission from those systems originate from the nuclei?'' and if not ``what are the spectral properties of any extended component?''  still remain unanswered.  This is where the superb sensitivity of space instruments, such as ISO, is essential.  We have found that in the ULIRGs studied here \\emph{more than $\\sim $\\,95\\% of the MIR emission seen by IRAS is confined within a few arcsecs of their central region}.  Obviously the relatively large pixel size of the ISOCAM detector places limitations in interpreting these findings. However, deconvolution tests of the central point source in each galaxy suggest that the corresponding nuclei are resolved and the physical diameter of the emitting region is contained within 1 to 2\\,kpc. Moreover, with the exception of the Superantennae where the MIR spectrum is dominated by the emission arising from the AGN of the southern galaxy, the bulk of the IR luminosity of IRAS\\,23128-5919 and IRAS\\,14348-1447 is powered by massive star formation. The fact that starbursts can dominate the MIR emission in galaxies with IR luminosities as high $\\sim$10$^{12}$\\,L$_{\\sun}$ had already been demonstrated in other ISOCAM-CVF \\citep{Tran2001} and ISO-SWS \\citep{Genzel1998} observations of ULIRGs, and is supported by our results.  Given that an active nucleus appears to be always present in the most energetic objects of the local Universe \\citep{Lutz1998}, our MIR data favor a luminosity threshold for the transition between starburst- and AGN-dominated galaxies which is higher than the IR luminosity of the galaxies in our sample. This is in agreement with the results of \\citet{Tran2001} who proposed that this transition takes place at L$_{\\rm IR}\\sim$10$^{12.5}$\\,L{$_{\\sun}$} and also found individual starbursts up to 10$^{12.65}$\\,L{$_{\\sun}$}. Our data also indicate is that such starbursts can be confined to the very central nuclear regions which may have important consequences in the probing how the instabilities fuel the inner regions of galaxies \\citep[e.g.][]{Combes2001}, as well as determining the nature of high redshift dusty sources \\citep[e.g.][]{Ivison2000}.  Another striking feature revealed in our observations is that in all three cases one galaxy seems to dominate the MIR energy output of the system by more than 75\\%. Could this be a record of the initial distribution of the amount of molecular gas available in each merging progenitor or could this suggest that in the later stages of interaction, the gas finally merges towards \\textit{one} component? If the latter were true one would expect that a sufficiently large quantity of gas could trigger and fuel both circumnuclear star forming activity and AGN-type activity at the core of a single object. This is evident in the southern galaxy of IRAS\\,19254-7245 which harbors an active nucleus as well as numerous massive star forming regions. As we mentioned in the introduction though the presence of a Seyfert nucleus is correlated with a MIR flux increase relative to the FIR luminosity of the entire galaxy, which is what one can actually derive from our observations when we compare the Superantennae with IRAS\\,14348-1447. IRAS\\,14348-1447 has indeed a much higher total IR luminosity despite its MIR flux being lower than that of the southern source of IRAS\\,19254-7245.  Furthermore, using the f$_{15\\mu m}$/H$\\alpha$ and f$_{15\\mu m}$/K ratios as probes of dust absorption and hot dust emission normalized to the mass of the galaxy, we find that in each interacting system it is always the most active galaxy of the system that exhibits the higher ratios. In each system, the most luminous galaxy contains a larger amount of molecular gas leading to the triggering/feeding of the starburst activity and/or an active nucleus.  Finally, we wish to stress once more that because of the limited spatial resolution in studying such distant sources, the diagnostics we have used in this paper address only the integrated MIR emission of each galaxy. Our difficulties to identify whether an active nucleus is solely responsible for the increase in the MIR luminosity relative to the FIR emission will not be resolved unless we can either clearly map the extent of the emitting region or obtain MIR spectra using very narrow slits. The upcoming launch of SIRTF which, despite the fact it has comparable spatial resolution to ISO, is equipped with a new generation of detectors of smaller pixel size, and in particular the use of its infrared spectrograph will help us improve upon our current results and provide conclusive answers to the issues which still remain unresolved to date."
    },
    "0206/astro-ph0206199_arXiv.txt": {
        "abstract": "Using data from the Two Micron All-Sky Survey (2MASS), we identify a population of infrared carbon (IR~C) stars with \\jk\\ $\\geq$ 2 in the Milky Way.  These stars are shown to trace the stellar bar previously identified in IR and optical surveys. The properties of C~stars strongly suggest that they are of intermediate age.  We conclude that the bar is likely to have formed more recently than 3~Gyr ago, and must be younger than 6~Gyr.  Implications and further tests of this conclusion are briefly discussed. ",
        "introduction": "It is by now well-established that there is a strong stellar bar in the inner disk of the Milky Way. A central bar was first hypothesized by \\citet{dev64}, although the observational evidence has only become overwhelming during the past decade (e.g., Kuijken 1996). Near and far-infrared source count maps have led numerous authors to identify a triaxial bar roughly 3--5 kpc long, with its near side in the first quadrant of the Galaxy \\citep[e.g.,][]{bli91,nak91,wei92,wei94}.  The OGLE survey extended these maps of the inner Galaxy to optical wavelengths and also found the bar morphology \\citep{sta94,sta97}.  Recent work based on 2MASS data \\citep{skr01} shows the full bar quite clearly.  Dynamical arguments suggest that bars in galaxies may be triggered by satellites and companions, intrinsic halo asymmetry, or as a result of disk instability.  Once formed, bars may be important for driving the evolution of galaxies through global angular momentum redistribution and increased rates of gas transport to galaxy centers.  As a result, barredness may affect a galaxy's star-formation history and nuclear activity. The dynamics of bars together with knowledge of bar ages may be crucial to understanding disk-halo interactions and merger histories. Despite the wealth of data on bar morphologies, little is known of their ages, except for the statistical result that bar frequency appears to decline with redshift for $z \\gtrsim 0.5$ \\citep{abr99}. If the stellar populations of the Milky Way bar can be age-dated, then the approximate time of the triggering event can be established.  This is a first step towards reconstructing the dynamical history of disk asymmetries in the Galaxy.  We know that infrared carbon (IR~C) stars reliably map out the Milky Way bar \\citep[][this paper]{skr01}. In \\S2, we briefly recapitulate the selection criteria for \\irc s that allow us to identify this population and to invert its magnitude distribution into a distance distribution.  \\S3 describes the arguments for an intermediate age for the \\irc s.  We argue that they are certainly younger than 6~Gyr, and probably younger than 3~Gyr, and therefore the bar must have formed more recently than these limits.  In \\S4, we discuss the implications of this result for studies of bar formation and suggest further work to refine our conclusion.  In short, if the Milky Way is typical, bars might be temporary features that can be successively re-excited during a galaxy's lifetime. ",
        "conclusions": "We have used 2MASS data to trace the structure of the Galactic disk and bar using the \\citet{skr01} sample of \\irc s with (\\jk ) $>$ 2. The epoch of strong star formation along galactic bars is expected to be brief \\citep[$\\lesssim$1~Gyr---][]{mar95,mar97}, and to occur during their formation.  Little or no star formation is expected to occur within the bar once it has become well-established, and stars that subsequently form in the disk cannot themselves join the bar.  Some support for this view comes from the observation that H {\\small II} regions are common in the bars only of late-type, active, or morphologically disturbed galaxies \\citep[e.g.,][and references therein]{mar97}. Therefore, the progenitors of the \\irc s were born prior to or during the formation of the bar, and hence their lifetimes give an upper limit to the bar age.  The oldest attested C~stars (in SMC clusters with [Fe/H] $\\approx -1.3$) are 5--7~Gyr old.   If all the C~stars in the Galactic bar were similar to the metal-poor SMC stars, reddened into our color window, the bar could thus be as old as roughly 6~Gyr.  Among the oldest star clusters with \\irc s are NGC~2121 (LMC, 3.2~Gyr), and {\\it possibly} Tr~5 (Milky Way, 2.8~Gyr).  M32 appears not to contain \\irc s \\citep{dav00}, and contains a large stellar population aged 3--5 Gyr \\citep{del01}.  We thus find it highly probable that the Milky Way bar is younger than 3~Gyr.  This time frame is intriguing, given the reports of a declining bar frequency among galaxies with redshifts $z \\gtrsim 0.5$ \\citep{abr99}, corresponding to lookback times of 5--6~Gyr.  Few age estimates have been made for the Milky Way bar.  \\citet{ng96} ascribed an 8--9~Gyr stellar population in Baade's Window to the bar, but the population's spatial distribution is not known, making the bar identification tentative.  \\citet{sev99} inferred an age of 7.5~Gyr from OH/IR stars with M $\\approx$ 1.3~M$_{\\sun}$; updated theoretical models \\citep{gir00} give 4.7~Gyr for this mass\\footnote{Because the OH/IR stars yield a similar age to the \\irc s, even a catastrophic misjudgment of the C~star fraction in Figure~1b does not invalidate our bar age limit.}.  The main-sequence turnoff of a 3~Gyr-old population should be readily traceable along the Galactic bar from V $\\approx$17 at the near end to V $\\approx$19 at the far end (modulo reddening differences).  Could the \\irc s be a trace population that wandered into the bar by chance? The LMC has M$_{\\mathrm{disk}}$ $\\approx$ 10$^{10}$ M$_{\\sun}$, and contains 2100 \\irc s.  \\citet{sta97} give M$_{\\mathrm{bar}}$ $\\approx$ 2 $\\times$ 10$^{10}$ M$_{\\sun}$, and we find 5300 sources with \\jk\\ $>$ 2 inside the bar radius.  For reasonable rates of contamination by OH/IR stars, the bar seems to have an \\irc\\ specific frequency comparable to the LMC's, arguing against the idea that 2MASS is seeing a trace young population within a much older bar.  2MASS observations of the LMC show that our color cut excludes roughly 75\\% of the N-type C~stars \\citep{nik00}.  If more of the C~stars could be traced, an accurate estimate of the fraction of intermediate-age stars in the bar could be made \\citep{aar85}.  Could the \\irc s be older than roughly 3~Gyr? Membership studies for Tr~5 and other potential C~star-bearing open clusters would help to empirically set the upper age limit for C~star formation. Further study of M32's AGB, to definitively measure its red extent, could push the likely age of the bar upwards, if its \\jk\\ reaches $\\approx 2.5$ \\citep{fre92}.  Bar dust lanes can have high molecular gas content, yet low star-formation rates \\citep[e.g.,][]{dow96}. Even in galaxies undergoing transient bursts (e.g., NGC~7479), the intensity of star formation along the bars is generally suppressed relative to that of the disk \\citep{lai99}. This suggests that a bar contains a snapshot of the disk stellar population at the time of its formation. However, barred galaxies typically have enhanced star formation at the ends of their bars and in their nuclei.  Could these young stars join the bar as they form?  A star can only be trapped in the bar by losing angular momentum as its orbit passes through a resonance.  This can only happen efficiently if the bar is actively evolving.  The existence of a well-described molecular ring in the Galaxy, presumably near corotation, suggests that our bar has a stable pattern speed, and therefore is not rapidly swallowing its own ends.  Bars are ubiquitous in numerical simulations; detailed studies by many authors show that the quasistatic gas response agrees well with observed morphology \\citep[e.g.,][]{ath92}.  This morphological coincidence suggests that the bars are at least several rotation times ($\\sim$1~Gyr) old, and that their pattern speeds are not rapidly evolving. On the other hand, a concentrated dark matter halo can cause a bar to lose angular momentum through dynamical friction.  \\citet{deb00} have argued that bars such as the Milky Way's are strongly braked by this effect, which would imply rapid bar evolution. However, \\citet{wei02} have hypothesized that a primordial bar that forms during the epoch of disk assembly will torque up the inner halo as it is braked, producing a core in the halo density distribution.  This initial, transient bar paves the way for a long-lived, stellar bar, since the altered halo mass profile does not slow bars as efficiently. The genesis of the current Milky Way bar remains to be explained. Tidal triggering by the LMC, the Sagittarius dwarf, or a now-merged satellite is a plausible origin \\citep{mur98,wei98,ves00}."
    },
    "0206/astro-ph0206220_arXiv.txt": {
        "abstract": "We report the discovery of a pulsar with period $P = 136$\\,ms and dispersion measure 308\\,cm$^{-3}$\\,pc in a deep observation of the supernova remnant (SNR) \\snr\\ with the Arecibo radio telescope.  Timing measurements of the new pulsar, J1930+1852, reveal a characteristic age of $P/2\\dot P = 2900$\\,yr and spin-down luminosity $\\dot E = 1.2 \\times 10^{37}$\\,erg\\,s$^{-1}$.  We have subsequently searched archival \\asca\\ X-ray data of this SNR, and detect pulsations with a consistent period.  These findings ensure that \\psr\\ is the pulsar that powers the ``Crab-like'' SNR~\\snr.  Together with existing \\chandra\\ observations of the SNR, we derive an X-ray pulsed fraction (2--10\\,keV) of $\\approx 27\\%$.  We also find that the cooling efficiency of the pulsar wind nebula (PWN) is intermediate between those of the Vela and Crab PWNe: $L_{X} \\mbox{(2--10\\,keV)} \\sim 0.002 \\dot E$.  \\psr\\ is a very weak radio source, with period-averaged flux density at 1180\\,MHz of $60\\,\\mu$Jy.  For a distance of 5\\,kpc, its luminosity, $\\sim 1$\\,mJy\\,kpc$^2$, is among the lowest for known young pulsars. ",
        "introduction": "\\label{sec:intro}  The supernova remnant (SNR) \\snr\\ is a close analogue of the Crab Nebula in several respects.  At radio wavelengths it has a filled-center morphology, a flat spectrum, and shows significant polarization \\cite{rfa+85,vb88}.  Its X-ray spectrum is non-thermal, and the X-ray and radio extents are comparable ($\\la 2'$; Lu, Aschenbach, \\& Song 2001)\\nocite{las01}.  It is a classic compact synchrotron nebula powered by its putative energetic pulsar, with no evidence for a thermal component corresponding to shocked ISM or stellar ejecta.  Using the \\chandra\\ X-ray Observatory, Lu et al.~(2002)\\nocite{lwa+02} have recently identified the central compact source (pulsar candidate) and imaged with $\\sim$ arcsecond resolution beautiful coherent structures (e.g., a ring) that are a manifestation of the relativistic pulsar wind interacting with the ambient medium.  The morphological and spectral properties of \\snr\\ revealed by \\chandra\\ leave no room for doubt as to the presence of a central pulsar.  It is nevertheless of considerable interest to detect actual pulsations, and to measure the period of rotation $P$ and its derivative $\\dot P$.  Because the SNR is powered entirely by the pulsar, knowledge of the input energy source (pulsar spin-down) luminosity ($\\dot E = 4 \\pi^2 I \\dot P/P^3$, where $I \\equiv 10^{45}$\\,g\\,cm$^2$ is the neutron star moment of inertia) would improve our understanding of the energetics of the nebula.  Also, the pulsar characteristic age $\\tau_c = P/2 \\dot P$ is a useful (if sometimes crude) measure of the SNR age, especially in cases like \\snr\\ where no independent reliable age estimate exists.  Finding radio and/or X-ray pulsations is also important because a significant sample of well-studied SNRs and their pulsars provides a window into the distribution of initial spin periods and magnetic fields of neutron stars.  Finally, detection of pulsations addresses the poorly constrained luminosity distribution, ``beaming fraction'', and hence Galactic population of young neutron stars.  The SNR~\\snr\\ was searched for a radio pulsar by Gorham et al.~(1996)\\nocite{gra+96} with the ``pre-upgrade'' Arecibo telescope. At a frequency of 1408\\,MHz, their 30\\,min observation of a 40\\,MHz-wide band reached a claimed sensitivity of 300\\,$\\mu$Jy for a pulsar with period $P \\sim 0.1$\\,s of duty-cycle 10\\% and dispersion measure $\\mbox{DM} \\sim 300$\\,cm$^{-3}$\\,pc (based upon their assumptions, we find their limiting sensitivity to be $\\sim 100\\,\\mu$Jy).  Following the striking \\chandra\\ observations of \\snr\\ \\cite{lwa+02}, we attempted the most sensitive search currently possible at Arecibo.  In this Letter we report the discovery of a 136\\,ms radio pulsar which, through a subsequent detection in archival \\asca\\ X-ray data, is confirmed to be the neutron star powering SNR~\\snr. ",
        "conclusions": "\\label{sec:disc}  The $P$ and $\\dot P$ measured for \\psr\\ imply a large spin-down luminosity $\\dot E = 1.2 \\times 10^{37}$\\,erg\\,s$^{-1}$, small characteristic age $\\tau_c = 2900$\\,yr, and surface magnetic dipole field strength $B = 3.2\\times10^{19} (P \\dot P)^{1/2} = 1.0 \\times 10^{13}$\\,G (Table~\\ref{tab:parms}).  These parameters are virtually identical to those of PSR~J1124$-$5916 in SNR~G292.0+1.8, and place it in the group of $\\sim 10$ pulsars with the highest values of $\\dot E$ and smallest apparent ages known, all of which are associated with SNRs (see Camilo et al.~2002a)\\nocite{cmg+02}.  However its parameters are substantially different from those of the Crab pulsar, which remains the only Galactic neutron star known with $P<50$\\,ms and $\\dot E > 10^{38}$\\,erg\\,s$^{-1}$.  The pulsar's characteristic age allows for a useful estimate of its actual age (and that of the SNR), given by $\\tau = 2 \\tau_c [ 1 - (P_0/P)^{n-1} ]/(n-1)$ under the assumption of constant magnetic moment, where $P_0$ is the initial period and $n$ is the braking index of rotation \\cite{mt77}.  For magnetic dipole braking, $n = 3$, and measured values span $2 \\la n \\la 3$ (see, e.g., Camilo et al.~2000)\\nocite{ckl+00}.  Initial periods of rotation are thought to range from $\\approx 10$\\,ms, to perhaps $\\ga 90$\\,ms (for PSR~J1124$-$5916; Camilo et al.~2002a)\\nocite{cmg+02}.  Considering the extremes of $n \\approx 2$ and $P_0 \\approx 90$\\,ms in turn, the age of J1930+1852/\\snr\\ likely lies in the range 1500--6000\\,yr.  The distance of SNR~\\snr\\ has been estimated from a measurement of the X-ray absorption column density, which is found to be about half the total Galactic absorption in this direction.  As the Galaxy here extends to $\\sim 10$\\,kpc from the Sun, Lu et al.~(2002)\\nocite{lwa+02} consider $d \\sim 5$\\,kpc.  Assuming the SNR to be associated with the star-forming region G53.9+0.3, Velusamy \\& Becker (1988)\\nocite{vb88} obtain $d \\sim 3.2$\\,kpc.  The dispersion measure of \\psr\\ (Table~\\ref{tab:parms}) provides an independent estimate: the Taylor \\& Cordes (1993)\\nocite{tc93} free electron density/distance model suggests $d \\sim 12$\\,kpc.  A newer model incorporating, among other improvements, discrete regions of enhanced ionized material (J.~M. Cordes \\& T.~J. Lazio 2002, in preparation), yields $d \\la 8$\\,kpc (T.~J. Lazio 2002, private communication), consistent with the estimate derived from X-ray measurements.  While the distance to this SNR/pulsar pair remains uncertain, we here retain $d \\sim 5$\\,kpc as a plausible estimate.  We now comment on some features of the pulsar wind nebula (PWN).  A composite \\chandra\\ image of \\snr\\ (Fig.~\\ref{fig:cxo}) demonstrates the trend of X-ray spectral softening from the inner to the outer regions of the nebula (as first indicated in the analysis of Lu et al.~2002)\\nocite{lwa+02}.  This is likely caused by a combination of synchrotron cooling and adiabatic expansion of the shocked wind material.  The 2--10\\,keV luminosity of the PWN/pulsar combination is $L_X = 2.2 \\times 10^{34} (d/5\\,\\mbox{kpc})^2$\\,erg\\,s$^{-1} \\sim 0.002 \\dot E$.  This is a factor of $\\sim 2$ lower than that of the PWN surrounding PSR~J1124$-$5916, a pulsar with identical spin parameters to \\psr\\ \\cite{cmg+02}\\footnote{Likewise the radio luminosity ($10^7$--$10^{11}$\\,Hz) of \\snr, $L_R \\sim 1 \\times 10^{33} (d/5\\,\\mbox{kpc})^2$\\,erg\\,s$^{-1} \\sim 0.0001 \\dot E$ \\cite{vb88}, is half that of the PWN powered by PSR~J1124$-$5916.  However, in one significant respect these two pulsar/SNR systems are different: PSR~J1124$-$5916 and its PWN are part of the text-book composite SNR~G292.0+1.8, embedded in a large and bright shell of stellar ejecta \\cite{hsb+01,prh+02}, whereas to date no thermal emission has been detected surrounding \\snr.}.  By contrast, the cooling efficiencies ($L_X/\\dot E$) of the PWNe surrounding the Vela \\cite{pksg01} and J2229+6114 \\cite{hcg+01} pulsars, with similar $\\dot E$, are much lower: $L_X \\la 10^{-4} \\dot E$; while that of the Crab Nebula is $\\sim 25$ times higher (e.g., Helfand \\& Becker 1987)\\nocite{hb87}.  The extent to which these differences are due to different ambient media, energy spectra of the injected pulsar wind (including the time evolution thereof), and ages, among other variables, is not clear.  \\medskip \\epsfxsize=8.0truecm \\epsfbox{f3.eps} \\figcaption[f3.ps]{\\label{fig:cxo} \\chandra\\ ACIS-S3 image of SNR~\\snr, color-coded by energy: 1.0--2.0\\,keV (red), 2.0--3.5\\,keV (green), and 3.5--8.0\\,keV (blue). The X-ray images in individual bands were adaptively smoothed in identical manner with a Gaussian filter in order to achieve a count-to-noise ratio of $\\approx 6$ in the final 1--8\\,keV image. North is to the top and East is to the left. }  \\medskip  With the discovery of \\psr, four very young and energetic pulsars have been uncovered in the past year, all with luminosities of $\\sim 1$\\,mJy\\,kpc$^2$ in the 1400\\,MHz radio band (Halpern et al.~2001; Camilo et al.~2002a, b; this Letter)\\nocite{hcg+01,cmg+02,csl+02}. These are well below the luminosities of previously known young pulsars (see Camilo et al.~2002a)\\nocite{cmg+02}, and it is now clear that such pulsars can be extremely faint.  Determining their intrinsic luminosity distribution requires disentangling observed ``pseudo-luminosities'' from beam-averaged values.  For example, do the broad wings of emission in \\psr\\ visible in Figure~\\ref{fig:radioprof} indicate nearly aligned rotation and magnetic axes with an impact angle grazing the outer boundary of a radio beam with possibly large averaged luminosity? Studies of polarized emission may provide important clues to constrain the beaming geometry of individual pulsars.  Naturally, a large sample of young pulsars is also desirable, together with upper limits from the most sensitive searches possible of well-selected targets.  In this regard, it is significant that new and upgraded radio telescopes are now available.  The discovery of \\psr\\ is an excellent example, as it could only barely have been made with the Arecibo telescope prior to its upgrade in the late 1990s: in addition to the increased bandwidth, we can now sample a lower-frequency band (1175 vs. 1400\\,MHz), where pulsars are brighter and the telescope's gain is higher.  Planned improvements in system temperature and bandwidth at Arecibo, the new Green Bank Telescope, and the remarkably productive Parkes telescope, along with the continued availability of the superb \\chandra\\ X-ray Observatory, offer the prospect of yet further progress in studies of faint young neutron stars."
    },
    "0206/nucl-th0206072_arXiv.txt": {
        "abstract": "Direct neutron capture on $^{62}$Ni is calculated in the DWBA and the cross sections in the energy range relevant for s-process nucleosynthesis are given. It is confirmed that the thermal value of the capture cross section contains a subthreshold resonance contribution. Contrary to previous investigations it is found that the capture at higher energies is dominated by p-waves, thus leading to a considerably increased cross section at s-process energies and a modified energy dependence. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206293_arXiv.txt": {
        "abstract": "The effective optical depth in the Ly$\\alpha$ forest region of 1061 low-resolution QSO spectra drawn from the SDSS database decreases with decreasing redshift over the range $2.5\\le z\\le 4$.  Although the evolution is relatively smooth, $\\tau_{\\rm eff}\\propto (1+z)^{3.8\\pm 0.2}$, at $z\\sim 3.2$ the effective optical depth decreases suddenly, by about ten percent with respect to this smoother evolution. It climbs back to the original smooth scaling again by $z\\sim 2.9$. We describe two techniques, one of which is new, for quantifying this evolution which give consistent results. A variety of tests show that the feature is not likely to be a consequence of how the QSO sample was selected, nor the result of flux calibration or other systematic effects. Other authors have argued that, at this same epoch, the temperature of the IGM also shows a departure from an otherwise smooth decrease with time. These features in the evolution of the temperature and the optical depth are signatures of the reionization of \\Hep. ",
        "introduction": "\\label{Int}  The importance of resonant scattering by neutral hydrogen in the intergalactic medium (IGM) was described by Gunn \\& Peterson (1965), who used the lack of a strong absorption trough in the spectra of high-redshift quasars to set limits on the amount of dispersed \\H.  Lynds (1971) noted that in the spectra of distant quasars there are many absorption features blueward of the Ly$\\alpha$ emission line; he interpreted the absorption features as Ly$\\alpha$ lines produced by intervening material. The mean absorption in the Ly$\\alpha$ forest depends mainly on the gas density and the amplitude of the ionising background (Rauch et al. 1997; Rauch 1998). The absorption increases rapidly with increasing redshift $z$ (e.g., Schneider, Schmidt \\& Gunn 1991; Songaila \\& Cowie 2002).  The optical depth $\\tau$ for a gas consisting primarily of ionized hydrogen and singly ionized helium, which is at density $(1+\\delta)\\equiv \\rho/\\langle\\rho\\rangle$ relative to the background density $\\langle\\rho\\rangle$ and is in photo-ionization equilibrium at redshift $z$, is \\begin{equation} \\tau(z) \\approx 0.7\\,\\left({\\Omega_b h^2\\over 0.019}\\right)^2\\, \\left({\\Omega_m h^2\\over 0.3\\times 0.65^2}\\right)^{-1/2}\\, \\left({1+z\\over 4}\\right)^{4.5}\\,{T_4^{-0.7}\\over \\Gamma_{12}} \\, {(1-Y)\\over 0.76}{(1-Y/4)\\over 0.94}\\,(1+\\delta)^2 \\label{eq:tau} \\end{equation} where $\\Omega_b h^2$ is the baryon density, $H_0=100h$ km s$^{-1}$ Mpc$^{-1}$ is Hubble's constant, $\\Omega_m$ is the matter density, and $Y$ is the helium abundance by mass (e.g. Peebles 1993, \\S 23). The temperature of the gas is $T_4\\equiv T/10^4K$, and $\\Gamma_{12}\\equiv \\Gamma/10^{-12}$ s$^{-1}$ is the photo-ionization rate.  Equation~(\\ref{eq:tau}) suggests that $\\tau$ should evolve rapidly. Various authors (e.g., Jenkins \\& Ostriker 1991; Hernquist et al. 1996; Rauch et al. 1997) have noted that measurements of the mean transmission $\\langle {\\rm exp}(-\\tau)\\rangle$ and its evolution constrain the parameters in equation~(\\ref{eq:tau}), such as the ratio $(\\Omega_b h^2)^2/(\\Omega_m h^2)^{1/2}$ (Rauch 1998), and the evolution of $T_4^{-0.7}/\\Gamma_{12}$ (McDonald \\& Miralda-Escud\\'e 2001).  Equation~(\\ref{eq:tau}) shows that, after the reionization of \\H, the optical depth is expected to decrease smoothly with time, unless, for example, there is a sudden injection of energy into the IGM. For instance, if the temperature of the gas increases by a factor of two at some epoch, then equation~(\\ref{eq:tau}) suggests that the optical depth $\\tau$ in the Ly$\\alpha$ forest would decrease by a factor of $2^{-0.7}\\sim 0.6$.  There is some evidence of a factor of two change in the temperature of the IGM at $z\\sim 3 - 3.5$ (e.g. Schaye et al. 2001).  Reimers et al. (1997; also see Heap et al.\\ 2000; Kriss et al.\\ 2001) found  evidence for a sharp increase in the \\Hep\\ opacity around $z\\sim 3$, which they associated with \\Hep\\ reionization.  Songaila \\& Cowie (1996) and Songaila (1998) have argued that the observed evolution of \\Cfour/\\Sifour metal line ratios requires a sudden hardening of the ionizing background around $z\\sim 3$, which is consistent with \\Hep\\ reionization. Schaye et al. (2000) and Theuns et al. (2002a,b) showed that \\Hep\\ reionization at $z\\sim 3.5$ results in a jump of about a factor of two in the temperature of the IGM at the mean density, and found evidence for such a jump by studying the distribution of line-widths in the Ly$\\alpha$ forest. In addition, Schaye et al. (2000) and Ricotti, Gnedin \\& Shull (2000) found that the gas is close to isothermal at redshift $z\\sim 3$, indicating that a second reheating of the intergalactic medium took place at $z\\sim 3$. This too might be interpreted as evidence of the reionization of \\Hep. (Numerical simulations of the observational signatures of \\Hep\\ ionization are also presented in e.g., Meiksin 1994 and Croft et al. 1997.) However, Boxenberg (1998) and Kim, Cristiani, \\& D'Odorico (2002) found no change in \\Cfour/\\Sifour, and analyses by McDonald et al. (2001) and Zaldarriaga, Hui, \\& Tegmark (2001) did not find a significant temperature change at these redshifts. Thus, both from metal line ratios, and from measurements of line widths, there is some evidence for \\Hep\\ reionization at $z\\sim 3 - 3.5$, and that this event is associated with an increase in the temperature of the IGM, although the strength of the evidence is still being questioned.  If \\Hep\\ were ionized at $z\\sim 3 - 3.5$, and this caused the temperature of the IGM to increase by a factor of two, then our simple estimate of an associated sixty percent decrease in $\\tau$ is not quite right. For instance, it ignores the fact that the extra electron liberated by the ionization can increase the optical depth.  However, for $Y\\sim 0.24$, the increase in the electron density from the electron released by \\Hep\\ ionization can increase $\\tau$ only  by seven or eight percent.  Although this goes in the opposite direction to the effect of the temperature increase, it is a substantially smaller effect.  Other important factors, which the simple sixty-percent estimate ignores, include the facts that the temperature change may be accompanied by a change in the temperature--density relation of the gas; that saturated lines which contribute to the optical depth will not be as strongly affected by a temperature change; and that a temperature increase may expand the gas, thus affecting peculiar velocities and complicating the relationship between temperature, line profile and optical depth. Nevertheless, the discussion above indicates that a sudden change in the temperature of the IGM may well be accompanied by a sudden change in the optical depth, although a precise estimate of the magnitude of the effect requires hydrodynamical simulations.  A sudden change in $\\tau$ means that the ratio of the mean absorption in the Ly$\\alpha$ forest to that in the spectrum of the quasar (hereafter QSO) should also change abruptly at the same time. That is, the quantity defined by Oke \\& Korycansky (1982), \\begin{equation} D_A\\equiv 1 - \\bar F, \\qquad {\\rm where}\\qquad \\bar F\\equiv {F_\\lambda({\\rm observed})\\over F_\\lambda({\\rm continuum})} \\equiv \\exp(-\\tau_{\\rm eff}) \\label{eq:DA} \\end{equation} should show a feature at $z\\sim 3 - 3.5$ if \\Hep\\ was ionized at that time. An advantage of studying the mean absorption, $D_A$, or transmission, $\\bar F$, is that it can be measured even in low resolution spectra for which individual line measurements are not possible. It is conventional to use the mean transmission to define an effective optical depth:  $\\tau_{\\rm eff}\\equiv -\\ln \\bar F$. Schneider, Schmidt \\& Gunn (1991) show that the mean transmission evolves significantly over the range $0<z<4.5$ (also see Press, Rybicki \\& Schneider 1993). The main goal of this paper is to see if this evolution is smooth, or has a feature in it.  For example, we would like to see if there is any evidence of a sudden drop in the effective optical depth in the Ly$\\alpha$ forest at $z\\sim 3 - 3.5$.  Section~\\ref{selection} describes how we selected our sample of $\\sim 10^3$ QSOs from the Sloan Digital Sky Survey (SDSS) database. The SDSS QSO selection algorithm itself is studied in some detail in Appendix~\\ref{seffects}.  We will be searching for a feature in the evolution of the mean transmission; this requires an accurate determination of the underlying intrinsic QSO spectrum.  This is the subject of Section~\\ref{contfit}.  We use two methods to do this. one method, which is a direct descendent of the one first used by Oke \\& Korycansky (1982), is described in Appendix~\\ref{method2}. The other is new; it exploits the fact that the continuum is a function of restframe wavelength, whereas the Ly$\\alpha$ effective optical depth is a function of observed wavelength. Both methods yield consistent results---the inferred continuum between the Ly$\\alpha$ and Ly$\\beta$ emission lines is not featureless, and, at $z\\sim 2.9-3.3$, there appears to be a feature in the otherwise smooth evolution of the mean transmission, and hence of $\\tau_{\\rm eff}$. Possible systematic effects which might affect our measurement are discussed in Section~\\ref{tauz} and in Appendix~\\ref{seffects}. A final section summarizes our findings.  Appendix~\\ref{app:skew} is somewhat tangential to the main subject of this paper:  it is a short demonstration of some effects which arise from the fact that the distribution of flux decrements in the Ly$\\alpha$ forest is highly non-Gaussian.  A comparison of this measurement with predictions from hydrodynamical simulations shows that our measurements can be interpreted as evidence for \\Hep\\ reionization at $z\\sim 3.2$ (Theuns et al. 2002). The implications for the evolution of the temperature of the IGM and the photo-ionization rate $\\Gamma$ will be presented in a future paper. ",
        "conclusions": "When applied to a sample of 1061 QSO spectra drawn from the SDSS database, two methods---which solve simultaneously for the shape of the QSO continuum and for the evolution of the mean transmission---give consistent results. Both methods show that the continuum in the wavelength range between the Ly$\\alpha$ and Ly$\\beta$ emission lines is not smooth, but has features in it.  The two methods also show that although the effective Ly$\\alpha$ optical depth decreases smoothly with time, $\\tau_{\\rm eff}\\propto (1+z)^{3.8\\pm 0.2}$, it drops by about 10 percent from $z\\sim 3.3$ to $z\\sim 3.1$, and it recovers to the original smooth scaling by $z\\sim 2.9$.  \\begin{figure} \\plotone{bernardi.fig13.ps} \\caption{Comparison of measurements of the evolution of the effective optical depth in the Ly$\\alpha$ forest.  Stars, diamonds, squares and small filled circles show measurements from 42 low resolution spectra by Sargent, Steidel \\& Bocksenberg (1989), 33 from Schneider, Schmidt \\& Gunn (1991), 42 from Zuo \\& Lu (1993), and the subset of 796 QSOs in the SDSS sample which had $S/N>4$ and were studied in this paper. Triangles and large filled circles show measurements in $\\sim 10$ higher resolution spectra by McDonald et al. (2000) and Schaye et al. (2000). Dotted line shows the evolution reported by Press, Rybicki \\& Schneider (1993), and dashed line shows the evolution given in Table~1.} \\label{taulit} \\end{figure}  A comparison of our measurement of $\\tau_{\\rm eff}$ with the findings of other authors is shown in Figure~\\ref{taulit}.  Stars, diamonds, squares and small filled circles show measurements from low resolution spectra of 42 QSOs by Sargent, Steidel \\& Bocksenberg (1989), 33 QSOs from Schneider, Schmidt \\& Gunn (1991), 42 QSOs from Zuo \\& Lu (1993), and 796 QSOs from the SDSS sample studied in this paper (the 796 spectra with $S/N>4$ out of the full sample of 1061 QSOs; the Ly$\\alpha$ forest was defined to span the range $1080-1160$~\\AA). Dotted line shows the evolution in the Schneider, Schmidt \\& Gunn sample reported by Press, Rybicki \\& Schneider (1993), and dashed line shows the evolution given in Table~1.  Large filled circles and open triangles show measurements from high resolution high signal-to-noise spectra by Schaye et al. (2000) and McDonald et al. (2000)---but note that although the two sets of analyses differed (McDonald et al. quote results in coarser redshift bins than do Schaye et al.), they were performed on essentially the same set of $\\sim 10$ QSO spectra.  Figure~\\ref{taulit} shows that our measurements are in general agreement with previous work based on low resolution, low signal-to-noise spectra (compare small filled circles with stars and diamonds). However, because our sample is so much larger than any previously available, it is possible to measure the evolution of $\\tau_{\\rm eff}$ in smaller redshift bins than was possible previously.  The figure also shows that our measurement of $\\tau_{\\rm eff}$ results in about ten percent less transmitted flux than suggested by recent measurements from higher resolution spectra (triangles and large circles).   At higher redshifts, where the absorption is large, it becomes increasingly difficult to estimate the continuum reliably for high resolution spectra. This, together with small number statistics may account for some of the discrepancy at $z>3.5$.  At lower redshifts, some of the discrepancy may arise because damped Ly$\\alpha$ systems and/or metal lines, which become increasingly abundant at low redshifts, have been removed from the higher resolution spectra, but are still present and contributing to the effective optical depth in our sample.  Although the mean transmission measured in noisy and low resolution spectra may yield a biased measure of the slope and amplitude of the evolution of the true effective optical depth (e.g., Steidel \\& Sargent 1987), it is difficult to see why this bias should lead to a {\\it feature} in $\\tau_{\\rm eff}(z)$. Thus, whereas the slope and amplitude of the evolution we find should be calibrated against simulations and other measurements, we feel we have strong evidence that the effective optical depth of the IGM changed suddenly around $z\\sim 3.2$.  The gradual evolution of the effective optical depth, and the strength of the feature superposed on it, both have implications for how the temperature and the photo-ionization rate evolve (equation~\\ref{eq:tau}). It is interesting that the feature we see in $\\tau_{\\rm eff}$ occurs at the same redshift range as the factor of two increase in temperature that Schaye et al. (2000) detected. Schaye et al. interpreted their measurement as evidence that \\Hep\\ reionized at $z\\sim 3.5$.  Hydrodynamical simulations show that our measurement of the evolution of $\\tau_{\\rm eff}$ is consistent with this interpretation (Theuns et al. 2002). The simulations can also help us understand if the mean scaling $\\tau_{\\rm eff}\\propto (1+z)^{3.8\\pm 0.2}$ we see leads to reasonable values for the amplitude and evolution of $\\Gamma$. It would be a significant accomplishment if the simulations were also able to reproduce the evolution of the skewed distribution around the mean---the latter being quantified by how the median optical depth depends on smoothing scale and on redshift (Figure~\\ref{fig:MBtaumedian}).  This is the subject of ongoing work.  The reionization of \\Hep\\ is expected to be proceed more gradually than for \\ion{H}{1}.  The fact that the feature appears relatively gradually in our data can be used to place constraints on how patchy the onset of reionization was.  When the SDSS survey is complete, it will be possible to compile a data set which is large enough to study different portions of the sky separately.  This will provide an even more direct constraint on the homogeneity of the Universe at the epoch of \\Hep\\ reionization.   \\bigskip  {\\em Acknowledgments}  This project started when Tom Theuns visited Fermilab in September 2001. We would like to thank him for a discussion which led to this measurement, and for helpful correspondence since then.   We also thank Paul Hewett, Patrick McDonald, Celine Peroux, and Uros Seljak.  Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, the University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington.  \\appendix"
    },
    "0206/astro-ph0206400_arXiv.txt": {
        "abstract": "The Nobeyama Millimeter Array (NMA) has been used to make aperture synthesis CO(1$-$0) observations of the post-starburst galaxy NGC 5195. CO(1$-$0) and HCN(1$-$0) observations of NGC 5195 using the Nobeyama 45 m telescope are also presented. High-resolution (\\timeform{1\".9} $\\times$ \\timeform{1\".8} or 86 pc $\\times$ 81 pc resolution at $D$ = 9.3 Mpc) NMA maps show a strong concentration of CO emission toward the central a few $\\times$ 100 pc region of NGC 5195, despite the fact that the current massive star formation is suppressed there. The face-on gas surface density, $\\Sigma_{\\rm gas}$, within the $r<2''$ or 90 pc region reaches $3.7 \\times 10^3$ $M_\\odot$ pc$^{-2}$ if a Galactic $N_{\\rm H_2}/I_{\\rm CO}$ conversion factor is applied. The extent of the central CO peak is about $5''$, or 230 pc, and is elongated along the E--W direction with two-armed spiral-like structures, which are typical for barred disk galaxies. The HCN-to-CO integrated intensity ratio on the brightness temperature scale, $R_{\\rm HCN/CO}$, is about 0.02 within the central $r < 400$ pc region. This $R_{\\rm HCN/CO}$ is smaller than those in starburst regions by a factor of $5-15$. These molecular-gas properties would explain why NGC 5195 is in a post-starburst phase; most of the {\\it dense} molecular cores (i.e., the very sites of massive star formation) have been consumed away by a past starburst event, and therefore a burst of massive star formation can no longer last, although a large amount of {\\it low density} gas still exists. We find a steep rise of the rotation velocity toward the center of NGC 5195. As a consequence, the critical gas surface density for a local gravitational instability of the gas disk becomes very high ($\\Sigma_{\\rm crit} \\sim 6.9 \\times 10^3$ $M_\\odot$ pc$^{-2}$), suggesting that the molecular gas in the central region of NGC 5195 is gravitationally {\\it stable}, in contrast to that of starburst galaxies. We propose that dense molecular gas can not be formed from remaining diffuse molecular gas because the molecular gas in the center of NGC 5195 is {\\it too stable} to form dense cores via gravitational instabilities of diffuse molecular gas. The deduced very high threshold density seems to be due to a high mass concentration in NGC 5195. The known trends on the occurrence and luminosity of nuclear star formation in early-type galaxies can be understood naturally if the high threshold density is characteristic for early-type galaxies. ",
        "introduction": "Recent studies on the molecular-gas properties in galaxies have revealed an intimate relationship between starbursts and high-density ($n_{\\rm H_2} > 10^4$ cm$^{-3}$) molecular gas; quantitative correlations between HCN(1$-$0) emission, which is an indicator of dense molecular gas contents, and star-formation tracers, such as a far-infrared (FIR) continuum and radio recombination lines, have been demonstrated (e.g., Solomon et al.\\ 1992; Zhao et al.\\ 1996; Paglione et al.\\ 1997). High-resolution images of HCN(1$-$0) emission with millimeter-wave arrays also show spatial coincidence of dense molecular material with the massive star-forming regions in starburst/star-forming galaxies (e.g., Paglione et al.\\ 1995; Kohno et al.\\ 1999, but see Aalto et al.\\ 2001).  Given the important role of dense molecular matter on massive star- formation, it would be then important to investigate how they are formed. A nonaxisymmetric distortion of the underlying potential in galaxies caused by various dynamical effects, such as disk instability and galaxy--galaxy interactions (e.g., Noguchi 1996), is often invoked as an efficient mechanism which stuff interstellar matter (ISM) on the large scale (from a few to a few $\\times$ 10 kpc) disk of galaxies into the very central (a few $\\times$ 100 pc) regions by removing the angular momentum of ISM (e.g., Shlosman et al.\\ 1989). To date, however, the physical process which could transform {\\it diffuse} ISM into {dense} molecular gas has not been addressed so often. Moreover, most of the observational studies on dense molecular gas have concentrated on luminous starburst galaxies, such as NGC 253 and M 82. It must be valuable to assess the distribution and physical properties of molecular gas in {\\it quiescent} and {\\it less star-forming} galaxies in order to seek a clue on dense gas formation.  In this paper, we report on high-resolution aperture synthesis CO(1$-$0) observations and simultaneous CO(1$-$0) as well as HCN(1$-$0) spectroscopy of the early-type barred galaxy NGC 5195, made with the Nobeyama Millimeter Array (NMA) and NRO 45 m telescope. As described below, NGC 5195 may be one of the ideal targets to investigate the molecular-gas properties in the galaxy whose star formation is currently at a very low level; NGC 5195 is a barred lenticular galaxy (SB0 pec; Sandage, Tammann 1981) at a distance of 9.3 Mpc (Tully 1988), known as a companion of the well-studied spiral galaxy M 51/NGC 5194. In spite of the abundant molecular gas in the central region of NGC 5195 (Sage, Wrobel 1989; Sage 1989, 1990; Aalto, Rydbeck 2001), the current massive star-formation seems to be suppressed there. Narrow-band imaging does not show any significant H$\\alpha$ {\\it emission}, and strong H$\\alpha$ {\\it absorption} dominates in the central $\\sim 10''$ region instead (Thronson et al.\\ 1991; Sauvage et al.\\ 1996; Greenawalt et al.\\ 1998; see also optical spectra by Filippenko, Sargent 1985). This Balmer absorption feature is attributed by a presence of numerous A-type stars, and few stars earlier than A4 exist in this region (Rieke 1988; Yamada, Tomita 1996). Considering the lifetime of A-type stars (from a few $\\times$ $10^8$ to $10^9$ yr) and the typical duration of starbursts (a few $\\times$ $10^7$; Thornley et al.\\ 2000 and references therein), it is strongly suggested that NGC 5195 experienced a nuclear starburst about $\\sim$ 1 Gyr before, and is now in the {\\it post-starburst phase}, where the OB stars produced by the starburst event have disappeared and only late-type stars (A stars and later) remain in this region. This situation is strikingly similar to the archetypical post-starburst galaxy NGC 4736 (Pritchet 1977; Rieke et al.\\ 1988; Walker et al.\\ 1988; Taniguchi et al.\\ 1996). A mid-infrared (MIR) study with ISOCAM by Boulade et al. (1996) has also revealed the post-starburst nature of NGC 5195. Consequently, the current star-formation rate (SFR) is very small; the H$\\alpha$ emission luminosity can be found by subtracting the underlying absorption feature by A stars, and is estimated to be $8.7\\times10^{37}$ erg s$^{-1}$ within a $2'' \\times 4''$ aperture (Ho et al.\\ 1997a). This is near the lowest end of $L$(H$\\alpha$) in their sample, containing about 500 nearby galaxies in the northern hemisphere. Note that optical spectroscopy suggests the nucleus of NGC 5195 may be a LINER (e.g., Ho et al.\\ 1997a), which could be a signature of a low-luminosity active galactic nucleus (AGN). Radio and MIR observations also show a compact and luminous source in the center of NGC 5195 (van der Hulst et al.\\ 1988; Boulade et al.\\ 1996). It is, however, unclear whether a compact non-stellar powered source exists in the nucleus of NGC 5195, because neither a compact UV core (Barth et al.\\ 1998) nor ionized neon lines in the MIR band (Boulade et al.\\ 1996), both are characteristic features of an AGN, has been detected. High-angular resolution X-ray observations of NGC 5195 also reveal an extended distribution of hot gas (Ehle et al.\\ 1995; Georgantopoulos et al.\\ 2002), in favour of a star-forming origin for the bulk of the X-ray emission. The radio and MIR concentrations, therefore, would be related to past starburst events (SNRs and hot dust heated by late-type stars). The parameters of NGC 5195 are listed in table 1. ",
        "conclusions": "The NMA observations of CO emission reveal that the post-starburst galaxy NGC 5195 does contain a significant amount of molecular gas, despite the fact that the current massive star formation is very inactive there; the molecular gas mass per unit area reaches about $\\Sigma_{\\rm gas} \\sim 3.7 \\times 10^3\\ M_\\odot$ pc$^{-2}$, which is comparable to those in nearby starburst galaxies. On the other hand, the star formation rate (SFR) per unit area, $\\Sigma_{\\rm SFR}$, which was calculated from $2'' \\times 4''$ aperture H$\\alpha$ data (Ho et al.\\ 1997a), is only about $\\Sigma_{\\rm SFR} \\sim 3.7 \\times 10^{-8}\\ M_\\odot$ yr$^{-1}$ pc$^{-2}$ (table 1). This SFR is smaller than those in typical starburst galaxies by an order of magnitude (e.g., Kennicutt 1998a). The resultant gas consumption time scale, \\begin{equation} \\tau_{\\rm gas} = \\frac{\\Sigma_{\\rm gas}}{\\Sigma_{\\rm SFR}}, \\end{equation} is about $\\sim 9 \\times 10^{10}$ yr, which is indeed comparable to $\\tau_{\\rm gas}$ in quiescent or normal galaxies, and significantly longer than those in infrared-luminous starburst galaxies (e.g., Kennicutt 1998a). In the following section we discuss the possible relationship between the physical properties of molecular gas and star formation in the center of NGC 5195.   \\subsection{Decrease of Dense Molecular Gas in the Nucleus of the Post-Starburst Galaxy NGC 5195}  We observed a very low $R_{\\rm HCN/CO}$ value of 0.02 in the center of NGC 5195. This is smaller than those in starburst galaxies by a factor of $5 - 15$. For instance, $R_{\\rm HCN/CO}$ in NGC 253 is reported to be in the range of $0.2 - 0.3$ (Helfer, Blitz 1993; Paglione et al.\\ 1995; Sorai et al. 2000), and starburst galaxies such as NGC 6946 and IC 342, which are less active compared with NGC 253 in terms of FIR luminosities, show $R_{\\rm HCN/CO}$ values of about 0.1 (Downes et al. 1992; Helfer, Blitz 1997). $R_{\\rm HCN/CO}$ values in the circumnuclear starburst regions of NGC 1068 (Helfer, Blitz 1995) and NGC 6951 (Kohno et al.\\ 1999) are also about 0.1 or so.  Because the critical gas density for collisionally excitation of HCN(1$-$0) emission ($n_{\\rm H_2} > 10^4$ cm$^{-3}$) is much higher than that of CO(1$-$0) ($n_{\\rm H_2} \\sim $ a few $\\times$ $10^2$ cm$^{-3}$), a comparison of the CO and HCN intensities is a measure of the gas density if both of the CO and HCN emission originate from the same volume. In observations of galaxies, the observing beams are often too large to resolve the individual cloud structures, and $R_{\\rm HCN/CO}$ could indicate the fraction of dense molecular gas to the total (including diffuse) molecular gas within the observing beam (Kohno et al.\\ 1999).  With this interpretation of $R_{\\rm HCN/CO}$ in galaxies, our data on NGC 5195 and its comparison with starburst galaxies would indicate that {\\it the mass fraction of dense molecular components to the total molecular gas}, which is traced by the $R_{\\rm HCN/CO}$ values, {\\it in the center of NGC 5195 is significantly small compared with that of starburst galaxies}. We suggest that this would be a reason why NGC 5195 is in a post-starburst phase despite the fact that it contains a large amount of molecular gas there; most of the dense molecular cores, which are the very sites of massive star formation, have been consumed away by starburst events which began about 1 Gyr before (see section 1). Also, a burst of massive star formation can no longer last, although a large amount of {\\it diffuse} molecular gas remains there.   We should note that the $R_{\\rm HCN/CO}$ values often depend on the aperture size of the observations, and attention must be paid to derive the intrinsic properties of molecular gas from the $R_{\\rm HCN/CO}$ values. (e.g., Helfer, Blitz 1993; Jackson et al.\\ 1996). For instance, $R_{\\rm HCN/CO}$ in the center of M 51 is 0.033 (Helfer, Blitz 1993) by the NRAO 12 m telescope ($55''$ for CO, and $71''$ for HCN), whereas it becomes about 0.08 (Sorai et al.\\ 2002) by the NRO 45 m telescope ($15''$ and $19''$ for CO and HCN, respectively). In higher resolution ($7''$) observations, it reaches about 0.3 or more (Kohno et al.\\ 1996). These effects are mostly due to the compactness of the HCN emitting regions, and therefore high-resolution observations are essential in order to address the relationship between dense molecular gas and massive star-forming regions. In the case of our NGC 5195 observations, the observing beams ($15''$ for CO and $19''$ for HCN) correspond to aperture areas of $r<340$ pc and $r<430$ pc for CO and HCN, respectively. Small $R_{\\rm HCN/CO}$ values of less than 0.05 have been reported in various galaxies (e.g., Helfer, Blitz 1993; Aalto et al. 1995), but are mostly observed with a much larger beam size ($\\sim 1'$) than ours, corresponding to a few $\\times$ kpc aperture area in their samples. We therefore suggest that the observed small $R_{\\rm HCN/CO}$ value in the central a few $\\times$ 100 pc of NGC 5195 is not the aperture size effect but reflects the intrinsic property of molecular gas, i.e., a decrease of dense molecular gas there. Our interpretation of a low $R_{\\rm HCN/CO}$ ratio seems to be supported by a low CO(2$-$1)/CO(1$-$0) ratio in NGC 5195; Sage (1990) observed CO(2$-$1) and CO(1$-$0) lines toward the center of NGC 5195 using the NRAO 12 m telescope, and found that the CO(2$-$1)/CO(1$-$0) ratio is as small as 0.5. They concluded that the molecular clouds in NGC 5195 are {\\it less dense} than those in typical spiral galaxies.    \\subsection{Gravitationally Stable Molecular Gas Disk in the Center of NGC 5195}  Given the close relationship between dense molecular gas and massive star formation/starburst, what kind of physical processes can play a major role in the formation of {\\it dense} molecular clouds from a {\\it low-density} molecular medium? One of the key processes could be gravitational instabilities of molecular gas; it is well known that SFR in the disk regions ($\\sim$ a few kpc scale) of galaxies can be described by means of the Schmidt law, $\\Sigma_{\\rm SFR} \\propto \\Sigma_{\\rm gas}^{1.4\\pm0.15}$, where $\\Sigma_{\\rm SFR}$ is the SFR per unit area, and $\\Sigma_{\\rm gas}$ is the gas mass per unit area (Kennicutt 1998a). Here it should be noted that the Schmidt law is valid only if the gas in the disk is gravitationally {\\it unstable}, i.e., the gas surface density, $\\Sigma_{\\rm gas}$, exceeds the critical gas surface density, $\\Sigma_{\\rm crit}$, given by \\begin{equation} \\Sigma_{\\rm crit} = \\alpha \\frac{\\sigma_v \\kappa}{\\pi G} = 73.9 \\times \\left(\\frac{\\alpha}{1}\\right) \\left(\\frac{\\sigma_v}{\\mbox{km s$^{- 1}$}}\\right) \\left(\\frac{\\kappa}{\\mbox{km s$^{-1}$ pc$^{-1}$}}\\right) M_\\odot \\mbox{\\ pc$^{-2}$}, \\end{equation} where $\\alpha$ is a dimensionless constant close to unity\\footnote{ In the case of a purely gaseous disk, $\\alpha$ is unity. Yet $\\alpha$ could be less than unity for a realistic two-fluid stability condition because the interaction between the stellar and gaseous disks could act to destabilize the gas disk (Jog, Solomon 1984). Martin and Kennicutt (2001) has shown that the best-fit value of $\\alpha$ is 0.69 $\\pm$ 0.2; we use this in the following calculation.}, $\\sigma_v$ is the velocity dispersion in the radial direction, and $\\kappa$ is the epicyclic frequency, \\begin{equation} \\kappa = \\left\\{ 2 \\frac{v(r)}{r} \\left[\\frac{v(r)}{r} + \\frac{dv(r)}{dr}\\right] \\right\\}^{0.5}, \\end{equation} where $r$ is the distance from the center and $v(r)$ is the rotation velocity at the radius of $r$ after correcting the inclination as 1/sin $i$. If the gas surface density, $\\Sigma_{\\rm gas}$, exceeds the threshold, $\\Sigma_{\\rm crit}$, then a uniform gas disk is unstable to form rings or clumps which can ultimately collapse into dense molecular gas fragments and form stars, eventually. This criterion, for local gravitational stability in thin isothermal disks, is often expressed as $Q = \\Sigma_{\\rm crit}/\\Sigma_{\\rm gas}$ (Toomre's $Q$ parameter; Toomre 1964); $Q>1$ ($Q<1$) means the gas disk is stable (unstable).  Surprisingly, this criterion has been successfully applied to much smaller scale disks, i.e., inner a few $\\times$ 100 pc regions of galaxies, to explain the star-formation properties of starbursts (e.g., NGC 3504, Kenney et al.\\ 1993) and rather normal (e.g., NGC 4414, Sakamoto 1996) galaxies. It is therefore natural to expect that the gravitational instabilities of molecular gas would be related to dense molecular gas formation. In fact, Kohno et al.\\ (1999) reported that $R_{\\rm HCN/CO}$ is enhanced in regions where the $Q$ parameter is almost equal to or smaller than unity, suggesting a connection between dense gas formation and gravitational instabilities of molecular gas.  In order to assess the stability of a molecular gas disk in the center of NGC 5195, we calculated $\\Sigma_{\\rm crit}$ in NGC 5195 using the above equations. The circular rotation velocities, obtained from the $pv$ map in figure 6, were employed to compute the epicyclic frequency, $\\kappa$. Note that our interest is just at the center of the gas disk, and that the gas motion can be treated as a rigid rotation in this region. In this case, the epicyclic frequency can be written just as \\begin{equation} \\kappa = \\sqrt{2} \\left[\\frac{v(r)}{r}\\right]. \\end{equation} From the $pv$ map in figure 6, we suppose that the largest circular rotation velocity occurs at about $r \\sim 1''$ or 50 pc and the $v(r)_{\\rm max}$ is about $100 \\mbox{\\ km s$^{-1}$}/{\\rm sin}\\ i = 160$ km s$^{-1}$, yielding a $\\kappa$ of 4.5 within $r<50$ pc. Note that $\\kappa$ must be a lower limit because the radius of $1''$ where the rotation rises to its peak could be an upper limit, given our limited angular resolution ($\\sim 2''$). Concerning the velocity dispersion of molecular gas, we here assumed an intrinsic gas velocity dispersion, $\\sigma_v$, of 30 km s$^{-1}$, which is similar to those in NGC 3504 (Kenney et al.\\ 1993). This is because a precise extraction of intrinsic gas velocity dispersion from the observed CO line width (figure 9) is not easy near the center of galaxies due to contamination of the steep velocity gradient within the observing beam. Consequently, we find a critical density, $\\Sigma_{\\rm crit}$, of $6.9\\times10^3\\ M_\\odot$ pc$^{-2}$, and that the $Q = \\Sigma_{\\rm crit}/\\Sigma_{\\rm gas}$ is larger than unity (1.9) in the center of NGC 5195. Hence, the molecular gas disk in the center of NGC 5195 would be gravitationally {\\it stable}; we suggest that this may be the reason why dense molecular gas is not formed from the large amount of diffuse molecular gas in NGC 5195, since they are {\\it too stable to form dense molecular clouds via gravitational instabilities of a diffuse molecular medium}.  Although there are some sources of error, we regard this conclusion as begin rather robust. Though one of the major uncertainties of the $Q$ parameter comes from $X_{\\rm CO}$, recent studies tend to show {\\it smaller} values than the $X_{\\rm CO}$ we adopted here by a factor of $2-3$ or more in the central regions of various galaxies (e.g., Nakai, Kuno 1995; \\cite{reg00}), including the Galactic Center (e.g., Oka et al.\\ 1998; Dahmen et al.\\ 1998). Even in the solar neighborhood, the latest extensive CO survey data of the Milky Way suggest a $X_{\\rm CO}$ of $1.8 \\times 10^{20}$ cm$^{-2}$ (K km s$^{-1}$)$^{-1}$ (Dame et al.\\ 2001), which is indeed smaller than the $X_{\\rm CO}$ adopted in our calculation. Thus the uncertainty of $X_{\\rm CO}$ does not weaken our conclusion. Although another source of error is the velocity dispersion, a large velocity dispersion of up to a few $\\times$ 10 km s$^{-1}$ can be observed even in rather quiescent galaxies; the Galactic Center is the case (e.g., Spergel, Blitz 1992). Given the finite spatial resolution, the epicycle frequency, $\\kappa$, could be a lower limit, indicating that the critical density can be even larger. It increases the $Q$ value then. All of these error sources thus indicate that the molecular gas in the center of NGC 5195 is indeed supercritical ($Q>1$) despite the large molecular gas mass there.  Supercritical ($Q>1$) molecular gas in post-starburst galaxies has also been reported at the centers of NGC 4736 (Shioya et al.\\ 1996) and NGC 7331 (Tosaki, Shioya 1997); both of them are prototypical post-starbursts (see Taniguchi et al.\\ 1996 and reference therein for NGC 4736, and see Ohyama, Taniguchi 1996 for NGC 7331). Gravitational instabilities of molecular gas may, therefore, be closely related to dense molecular gas formation and, in turn, massive star-formation in galaxies.   \\subsection{Star-Formation Properties in Early-Type Galaxies}  It should also be addressed why the molecular gas disk in NGC 5195 is supercritical. As mentioned above, the gas surface density, $\\Sigma_{\\rm gas}$, is rather comparable to those in nuclear starburst regions of galaxies. The major difference between starbursts and the post-starburst galaxy NGC 5195 may be the very high critical density for the gravitational instability; the deduced $\\Sigma_{\\rm crit}$ of $6.9\\times10^3\\ M_\\odot$ pc$^{-2}$ is significantly high compared with the central gas surface density in disk galaxies (e.g., Sakamoto et al.\\ 1999b).  This very high threshold may be caused by the large mass concentration in early-type galaxies (Kennicutt 1989). The large concentration of matter results in a steep rise of the rotation velocity, as observed in NGC 5195. This increases the threshold density, because $\\kappa$ becomes very high in this case. In other words, molecular gas at the centers of early-type galaxies tends to be stabilized due to a small gas mass-to-dynamical mass ratio, $M_{\\rm gas}/M_{\\rm dyn}$, even if a significant molecular gas mass concentration, comparable to late-type disk galaxies, occurs (e.g., Young, Scoville 1991).  These properties observed in the S0 galaxy NGC 5195 may explain the difference in the star-formation properties along the Hubble sequence. As reviewed by Kennicutt (1998b), it is now established that the detection frequency of nuclear star formation is lower in early-type galaxies than in late-type ones (e.g., the detection rate is about 8\\% in S0, whereas it increases about 50\\% in Sb and reaches 80\\% in Sc-Im; Ho et al.\\ 1997b), whereas the strength of nuclear star formation shows the opposite trend (i.e., the average extinction-corrected H$\\alpha$ luminosities in S0--Sbc galaxies is nine-times higher than in Sc galaxies; Ho et al.\\ 1997b). The high critical density due to the high mass concentration in early-type galaxies may explain these apparently incoherent trends; in early-type galaxies, the occurrence of nuclear star formation is low because nuclear star formation cannot occur until the gas surface density exceeds the higher threshold density compared with late-type galaxies. In order to exceed the threshold, it may take a longer time or many triggers to transport the larger amount of molecular gas into the centers of early-type galaxies. Once the gas surface density exceeds the threshold, however, the star-formation rate must be very high, because the gas surface density already reaches a very high value compared with late-type galaxies (Kennicutt 1989). Our high-resolution CO observations of the early type galaxy NGC 5195 do support this scenario. A CO survey of early-type galaxies using millimeter-wave interferometers and a comparison of the gas surface and the threshold densities with those in late-type spirals must be valuable to statistically test the picture described above regarding the difference in star formation along the Hubble sequence."
    },
    "0206/astro-ph0206414_arXiv.txt": {
        "abstract": "Air shower simulation programs are essential tools for the analysis of data from cosmic ray experiments and for planning the layout of new detectors. They are used to estimate the energy and mass of the primary particle. Unfortunately the model uncertainties translate directly into systematic errors in the energy and mass determination.  Aiming at energies $> 10^{19}$ eV, the models have to be extrapolated far beyond the energies available at accelerators.  On the other hand, hybrid measurement of ground particle densities and calorimetric shower energy, as will be provided by the Pierre Auger Observatory, will strongly constrain shower models.  While the main uncertainty of contemporary models comes from our poor knowledge of the (soft) hadronic interactions at high energies, also electromagnetic interactions, low-energy hadronic interactions and the particle transport influence details of the shower development.  We review here the physics processes and some of the computational techniques of air shower models presently used for highest energies, and discuss the properties and limitations of the models. ",
        "introduction": "The cosmic ray (CR) energy spectrum extends up to $3\\times 10^{20}$ eV. The presence of the highest-energy particles (Ultra High-Energy CRs, UHECRs) poses an enigma, since many good arguments suggest that they should not be observed.  This apparent contradiction has stimulated a variety of more exotic explanations of their existence. The enigma can only be solved by an experiment that can provide a much larger event statistics than the about 20 events with $E > 10^{20}$ eV measured in the past 35 years. Knowing the form of the energy spectrum of the CR particles, their arrival direction distribution over the whole sky, and possibly even their mass composition, would allow us to test some of the hypotheses about their origin and help to identify their sources. At present experimental results suggest that UHECRs are protons or nuclei, as for CRs at much lower energies.  However, many of the more exotic models of UHECR origin predict also photons and neutrinos.  The measurement of extensive air showers (EAS) is presently the only way to study CRs with energies above about 10$^{15}$ eV. The properties of primary cosmic rays have to be deduced from the development of the shower in the atmosphere and from the particle ratios in the shower. The incident direction can easily be reconstructed from the arrival times of shower particles at different positions at the observation level and the primary energy is approximately reflected in the total number of secondary particles produced.  The mass of the primary particle is more difficult to measure. It is reflected, in a subtle way, in the shower form, specifically the height of the shower maximum, and in the muon-to-electron ratio of the shower.  The Pierre Auger Observatory is conceived to measure CRs with energies $> 10^{19}$ eV with good statistics over the whole sky \\cite{auger}. It will consist of two detector sites, one in the southern and one in the northern hemisphere.  Each site covers an area of 3000 km$^2$ and combines two techniques to measure (i) the particle distribution at observation level with an array of water Cherenkov detectors and (ii) the longitudinal shower development via optical imaging of the fluorescence light in the atmosphere during clear moonless nights ($\\approx$ 10\\% of the total time). This hybrid detection provides a way to inter-calibrate both sub-systems and to control systematic uncertainties.  The energy determination via the fluorescence light is basically calorimetric and therefore much less model-dependent than the energy reconstruction from particle densities at ground level.  The southern site is presently under construction in Argentina. The size of the site was chosen to register about 5000 events with energies $> 10^{19}$ eV and about 40-80 events above 10$^{20}$ eV per year.  Since experiments at energies $> 10^{15}$ eV cannot be calibrated with a test beam the interpretation of EAS measurements is performed by comparing experimental data with model predictions of the shower development in the atmosphere. Therefore quantitative results rely on the model assumptions and on the quality of the simulation of particle interactions and transport in the atmosphere.  The detailed shower development is far too complex to be fully described by a simple analytical model. Therefore it is usually modeled by Monte Carlo (MC) simulation of transport and interaction of each individual shower particle, employing our present knowledge on interactions, decays and particle transport in matter.  While the electromagnetic interaction (responsible for electromagnetic sub-showers, ionization, Cherenkov light production, ...) and the weak interaction (responsible for decays of unstable secondaries) are well understood, the major uncertainties in EAS simulation arise from the hadronic interaction models.  With the present theoretical understanding of soft hadronic interactions, i.e. those with a small momentum transfer, one cannot calculate interaction cross-sections or particle production from first principles. Therefore, hadronic interaction models are usually a mixture of fundamental theoretical ideas and empirical parametrisations tuned to describe the experimental data at lower energies.  The large extrapolation (over 6 orders of magnitude in energy) needed from experimental accelerator data to CR interactions is the second major source of uncertainty, and with an uncertain interaction model it is difficult to determine the energy spectrum and the composition of CRs.  As a consequence of the better understanding of hadronic and nuclear interactions at high energies, and the increase in computing power, shower models have improved dramatically over recent years.  Also the understanding of the measuring process with a detector has markedly advanced. It is now possible to describe consistently experimental results over a wide range of energies and even to test hadronic models on the 20\\% level.  In this paper we review the physics and techniques of state-of-the-art modeling of extensive air showers, specifically those important for energies of $10^{19}$ eV or above.  In Sec. \\ref{sec-had} and \\ref{sec-em} we describe the modeling of hadronic and electromagnetic interactions, respectively.  Sec. \\ref{sec-neu} summarizes briefly the simulation of showers induced by photons and neutrinos or other more exotic primary particles.  Statistical thinning techniques that allow simulation of showers at the highest energies in a finite time are discussed in Sec.  \\ref{sec-thin}.  Then two shower simulation programs are described and compared in Sec. \\ref{sec-programs}. Some simulation results are shown in Sec. \\ref{sec-results}. ",
        "conclusions": "In the last decade much has changed in air shower simulations. Through better understanding of the relevant hadronic interactions and the massive increase in computer performance very elaborate Monte Carlo models became feasible that track individual particles through atmosphere and detectors and simulate all their interactions with matter in great detail.  Predictions of air shower programs have become much more quantitative and describe the experimental data sufficiently well for data analysis and the design of new experiments.  The weakest point in the simulations are the high-energy soft hadronic interactions, which are not well examined at accelerators so far, and the extrapolation to energies much beyond those available at accelerators.  The most successful models are based on the Gribov-Regge theory of multi-Pomeron exchange.  There is a clear trend of convergence between different hadronic interaction models, due to objective improvement of the physics input. However, a level of systematic uncertainty of about 20\\% is likely to persist, since many of the input data, on which the models are built, appear to have uncertainties of this size.  Once observed, larger inconsistencies between data and simulations offer a possibility to improve the model assumptions.  The calorimetric measurement of the shower energy via the fluorescence detectors in coincidence with the charged particle measurements in the Auger experiment will assist with stringent tests of the shower models and promises the  determination of an absolute energy scale as well as setting constraints on the shower models.  Within the Auger Collaboration the two programs CORSIKA and AIRES are used.  They account for all processes necessary for air showers simulations up to 10$^{21}$ eV.  The programs complement each other in the sense that CORSIKA uses more elaborate interaction models and AIRES is faster. CORSIKA and AIRES agree to better than 20\\% for the basic shower parameters observed by the Pierre Auger Observatory."
    },
    "0206/astro-ph0206308_arXiv.txt": {
        "abstract": "We study the absorption along lines of sight toward high-$z$ radio sources caused by the $21 \\cm$ transition of neutral hydrogen in the intergalactic medium (IGM) before reionization.  Using semi-analytic methods, we compute the number density of observable features caused by both ``minihalos'' (bound objects that are unable to cool efficiently because of their small virial temperatures) and protogalactic disks.  We show that both sets of features should be observable by the next generation of low-frequency radio telescopes, including the \\emph{Low Frequency Array} and the \\emph{Square Kilometer Array}, provided that sufficiently bright background sources exist.  The statistics of minihalo absorption features seen along lines of sight to radio-loud quasars offer a way to measure the evolution of the radiation background and the IGM temperature with cosmic time.  Intersections with disks are much less common but cause significantly deeper absorption features that would be visible in the spectra of both radio-loud quasars and gamma-ray bursts (GRBs).  The absorption feature caused by HI in the host galaxy of a GRB should be observable, offering a route to determine spectroscopically the burst redshift. ",
        "introduction": "\\label{intro}  Despite its apparent simplicity, effective observational probes of the intergalactic medium (IGM) at high redshifts are difficult to find. Between recombination at $z \\sim 1000$ and reionization at $z \\sim 6$, the IGM was almost entirely neutral.  In such a medium, the optical depth to Ly$\\alpha$ absorption at the frequency corresponding to a redshift $z$ is $\\tau_\\alpha \\approx 6.45 \\times 10^5 x_{\\rm HI} [(1+z)/10]^{3/2}$ \\citep{gp}, where $x_{\\rm HI}$ is the neutral fraction and where we have assumed the currently favored cosmological parameters (see below).  This enormous optical depth renders the Ly$\\alpha$ forest, which is the most important tool for studying the (highly ionized) IGM at low and moderate redshifts, almost useless for detailed studies of the pre-reionization IGM.  For example, the detection of a complete absorption trough shortward of the Ly$\\alpha$ resonance wavelength in the spectrum of the quasar SDSS 1030+0524 may be the first evidence for the end of the reionization era \\citep{becker}.  However, studies of rest-frame UV and optical quasar spectra are unable to place stronger constraints on the neutral fraction than $x_{\\rm HI} \\ga 10^{-2}$ \\citep{fan}.  Attempts to constrain other parameters, such as the thermal state of the IGM, the radiation background, or the density structure, are also compromised by the extremely large optical depth at the Ly$\\alpha$ resonance.  Clearly, we cannot obtain an in-depth understanding of the IGM before and during reionization through studies of the Ly$\\alpha$ forest.  An alternative is to observe a weaker transition at the other extreme of optical depth: the $21 \\cm$ hyperfine transition of neutral hydrogen. To date, there have been several theoretical studies of emission or absorption of the neutral IGM against the cosmic microwave background (CMB) in this transition \\citep{scott,kumar,mmr,tozzi,shaver,iliev}. However, the predicted signals are extremely weak.  Even the next generation of radio telescopes, such as the \\emph{Square Kilometer Array}\\footnote{See, e.g., http://www.usska.org/main.html.} (SKA), will only be able to detect exceptionally large objects (e.g., \\citealt{scott,kumar}).  Another strategy is to study the angular structure of the emission in order to constrain statistically the structure formation process at high redshift (e.g., \\citealt{tozzi,iliev}).  However, source confusion with low-$z$ faint radio sources \\citep{dimatteo} and free-free emission by the intergalactic medium (Loeb 1996) are likely to compromise such attempts.  Alternatively, \\citet{carilli} have pointed out that luminous high-redshift radio-loud quasars serve as ideal background sources against which absorption by intervening gas can be seen: with a bright background source it becomes much easier to identify the weak absorption features expected to be produced by the IGM.  Along such lines of sight one can map the ``21 cm forest'' of redshifted hyperfine absorption lines and hence study the neutral IGM over the redshift range $6 \\la z \\la 10$.  Using a simulation, \\citet{carilli} found that the high-$z$ analog of the Ly$\\alpha$ forest of filaments and other overdense regions produces recognizable absorption features in the spectrum of a radio-loud quasar.  They also demonstrated that such observations are feasible with an instrument similar to current designs for the SKA, provided that sufficiently bright radio-loud quasars exist at these early epochs.  Large scale numerical simulations, such as those studied by \\citet{carilli}, are unable to resolve collapsed gas clouds on the smallest mass scales, although these clouds represent the most abundant clumps in the high-redshift IGM.  In this paper we complement their study by exploring semi-analytically the statistics of these compact absorbing systems. Our semi-analytic approach allows us to examine the sensitivity of the results to changes in the input parameters and to identify the physical quantities that are best measured by radio absorption spectra.  As we demonstrate in the following sections, 21 cm absorption spectra probe the radiation background and thermal state of the IGM at high redshifts and therefore could be instrumental in testing models of structure formation and reionization.  The smallest bound objects in the IGM (commonly referred to as ``minihalos'') have virial temperatures below the threshold for atomic hydrogen line cooling, and so (in the absence of a substantial abundance of molecular hydrogen, H$_2$) they cannot cool and collapse to form protogalaxies.  Although minihalos cannot form stars, they are important for determining the mean clumping factor of the IGM and for screening ionizing radiation from other objects \\citep{barkana02}. Their properties are sensitive to the uncertain presence of molecular hydrogen; if H$_2$ forms in sufficient quantities it can act as a cooling channel for these minihalos (\\citealt{barkana01}, and references therein), strongly suppressing their number density (and hence increasing the global star formation rate).  Detection of these halos would provide important insights into the physics of hierarchical structure formation prior to reionization.  Halos with virial temperatures above the hydrogen line cooling threshold will collapse to form stars.  Thus we expect some fraction of all lines of sight to penetrate protogalactic disks.  Such absorption systems would be the pre-reionization analogs of damped Ly$\\alpha$ absorbers.  While disk intersections are rare, they can provide important information about the state of the neutral hydrogen in high-$z$ galaxies, including the distribution of disk masses, star formation rates, and gas temperature.  In \\S \\ref{model} we derive the optical depth to redshifted $21 \\cm$ radiation through both the diffuse IGM and collapsed objects.  We show our numerical results in \\S \\ref{results} and conclude in \\S \\ref{disc}.  Throughout the discussion we assume a flat, $\\Lambda$--dominated cosmology, with density parameters $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$, $\\Omega_b = 0.05$, in matter, cosmological constant, and baryons, respectively.  In the numerical calculations, we assume $h=0.7$, where the Hubble constant is $H_0 = 100 h \\hunits$. ",
        "conclusions": "\\label{disc}  We have shown that absorption at the 21 cm transition of neutral hydrogen in both minihalos and protogalactic disks at high-redshifts can cause non-negligible absorption in the low-frequency spectra of high-redshift radio sources.  Using a semi-analytic approach, we have found that the density of minihalo absorption features is primarily a function of the kinetic temperature of the IGM (which determines the minimum mass of the minihalos) and the radiation background (which couples the spin and kinetic temperatures of the hydrogen gas).  We predict that the average radiation background can be measured by the continuum absorption and the number density of very weak lines, while the IGM temperature can be measured by the density of narrow, relatively strong absorption lines.  Optical depths reach $\\tau_0 \\sim 0.02$ for objects common enough to appear along typical lines of sight.  Studies of absorption radio spectra therefore offer a probe of the era before reionization and can help constrain the history of early star formation.  We find, unsurprisingly, that disks are only rarely intersected but tend to have much higher optical depths.  The principal shortcoming of our semi-analytic approach is its inability to accurately describe the filamentary structure of the IGM and the spatial clustering of halos.  Using a simulation, \\citet{carilli} have also studied 21 cm absorption along lines of sight to high-redshift radio sources.  While the mass resolution of their simulation was too coarse to include the minihalo population, they were able to capture the geometry and structure of filaments. They find that both the typical optical depths and the number densities of filament absorption lines are comparable to or slightly larger than our estimates for those from minihalos.  An observed absorption spectrum will then be a mixture of lines from these two populations of absorbers.  The clustering of minihalo lines may yield information on bias or feedback from nearby cooled objects, but a proper study of such an effect will require numerical simulations or more detailed semi-analytic models (e.g., \\citealt{scann}).  Can these signals be detected by future radio telescopes?  First, note that the relevant frequency range is $100-200 \\mhz$, corresponding to $z \\sim 13$--$6$; thus, a low-frequency instrument is necessary.  The transmission curves in \\S 3.2 ignore noise.  In a real observation, the noise level will be determined by a combination of the telescope characteristics and the brightness of the background source.  Because the specifications of the next generation of telescopes are not yet settled, we have chosen not to present results specifically tailored to any particular instrument.  Instead, we can simply estimate the required source brightness, $S_{\\rm min}$, in order to observe a single absorption feature, given the telescope's effective area $A_{\\rm eff}$, system temperature $T_{\\rm sys}$, and channel width $\\Delta \\nu_{\\rm ch}$, for a specified signal-to-noise ratio $S/N$ and integration time $t$.  In the small $\\tau_0$ limit and under the assumption that $\\Delta \\nu_{\\rm ch} \\ll \\Delta \\nu_{\\rm obs}$, the result is \\bq S_{\\rm min} = 12 \\mjy \\, \\left( \\frac{S/N}{5} \\right) \\left( \\frac{0.02}{\\tau_0} \\right) \\left( \\frac{ 2 \\times 10^3 \\aefftsys}{A_{\\rm eff}/T_{\\rm sys}} \\right) \\left( \\frac{10 \\days}{t} \\right)^{1/2} \\left( \\frac{\\khz}{\\Delta \\nu_{\\rm ch}} \\right)^{1/2}. \\label{eq:sflux} \\eq Current designs for the SKA call for $A_{\\rm eff}/T_{\\rm sys} \\sim 2 \\times 10^3 \\aefftsys$, while the \\emph{Low Frequency Array}\\footnote{See http://www.lofar.org/index.html.} (LOFAR) is expected to have $A_{\\rm eff}/T_{\\rm sys} \\sim 4.5 \\times 10^2 \\aefftsys$.  A channel width of $1 \\khz$ should be achievable with LOFAR while maintaining a total spectral coverage $\\sim 4 \\mhz$ and with SKA while maintaining a total spectral coverage of $\\sim 10 \\mhz$.  Do objects of this brightness exist at sufficiently high redshifts? \\citet{haiman01} have shown that, in the context of hierarchical cold dark matter models, black holes with masses $\\ga 10^8 \\msun$ can plausibly exist even beyond $z \\sim 10$. While we currently have no observational constraints on this model, \\citet{carilli} have also presented a plausibility argument for the existence of luminous radio-loud quasars at $z \\ga 7$.  By extrapolating the radio galaxy luminosity function to high redshifts using the observed decline of optical quasars, they estimate $\\ga 2000$ sources on the sky of sufficient brightness to be useful.  Furthermore, they point out that a bright radio galaxy at $z=5.2$ has already been detected \\citep{breugel}, and there is no obvious environmental reason for a high-$z$ cutoff.  Because of their high optical depths, disks could be detected along lines of sight to much fainter sources.  In this case, equation (\\ref{eq:sflux}) must be modified to apply to large optical depths through the substitution $\\tau_0 \\rightarrow 1-\\exp(-\\tau_0)$.  For example, detecting a disk with $\\tau_0=1$ given the telescope parameters in equation (\\ref{eq:sflux}) would require an average source flux $\\sim 0.4 \\mjy$.  Extreme starburst galaxies and faint active galactic nuclei could satisfy this criterion and of course will be more common than the bright quasars useful for minihalo studies. However, the disk intersection probability is sufficiently small that only a small fraction of all lines of sight will pass through intervening disks, so searches using these sources are unlikely to be practical given the long integration times required.  Another particularly interesting type of source is a gamma-ray burst (GRB). In this case, the rarity of disks is not a limit because the host galaxy is necessarily present at the redshift of the burst.  (While the host galaxy of a faint AGN or a starburst may also be visible, the gas is much more likely to be ionized on large scales around such sources.)  The smooth power-law spectrum of GRBs is an ideal continuum for these observations.  The typical peak flux of a GRB is $\\la 1 \\mjy$ \\citep{ciardigrb}; however, the low-frequency spectra ($\\la 10 \\ghz$) of GRBs typically suffer from synchrotron self-absorption \\citep{piran}, reducing the flux at $1420 \\mhz$ by a factor of a few from the peak value (e.g., \\citealt{galama}).  GRBs are therefore too dim for the study of minihalo absorption features, but absorption by the host disk should be observable provided that $\\tau_0 \\ga 2$. Such a detection could confirm the redshift of the burst and provide information on the neutral hydrogen content of its host galaxy."
    },
    "0206/astro-ph0206078_arXiv.txt": {
        "abstract": "{ An extensive modelling of CO line emission from the circumstellar envelopes around a number of carbon stars is performed. By combining radio observations and infrared observations obtained by ISO the circumstellar envelope characteristics are probed over a large radial range.  In the radiative transfer analysis the observational data are consistently reproduced assuming a spherically symmetric and smooth wind expanding at a constant velocity.  The combined data set gives better determined envelope parameters, and puts constraints on the mass loss history of these carbon stars.  The importance of dust in the excitation of CO is addressed using a radiative transfer analysis of the observed continuum emission, and it is found to have only minor effects on the derived line intensities.  The analysis of the dust emission also puts further constraints on the mass loss rate history.  The stars presented here are not likely to have experienced any drastic long-term mass loss rate modulations, at least less than a factor of $\\sim$5, over the past thousands of years. Only three, out of nine, carbon stars were observed long enough by ISO to allow a detection of CO far-infrared rotational lines. ",
        "introduction": "Mass loss is associated with many phases of stellar evolution.  It is of particular importance during the final evolutionary stage of low to intermediate mass stars, the asymptotic giant branch (AGB), when it increases dramatically. Its overall characteristics, e.g., magnitude, geometry, and kinematics, are reasonably well established, but to what extent such effects as asymmetry, non-homogeneity, temporal variability, etc., are important is not known. During the AGB-phase, stars in the mass range about 1.5--4\\,M$_{\\sun}$ and with 'normal' chemical compositions start to evolve into carbon stars, i.e., stars with more carbon than oxygen in their atmospheres, through the dredge-up of freshly synthesized carbon in He-shell flashes \\citep{Straniero97, Busso99}. It appears that all AGB carbon stars are losing mass at a rate in excess of 10$^{-8}$\\,M$_{\\sun}$\\,yr$^{-1}$ \\citep{Olofsson93a, Schoeier01}.  Approximately 10\\% of all carbon stars appear to have been subject to drastic changes in their mass loss rates over relatively short periods of time (a few thousand years).  Most notable are the CO radio line observations of detached circumstellar envelopes \\citep{Olofsson96, Olofsson00, Lindqvist99, Schoeier01}.  The time interval between these, supposedly reoccuring, mass ejections is estimated to be about 10$^5$\\,years, and they are possibly linked with the He-shell flashes predicted to occur in carbon stars \\citep{Olofsson90, Schroeder99, Steffen00}. Less dramatic mass loss rate modulations have recently been observed by \\citet{Mauron99, Mauron00} in the form of a multiple-shell structure around the high mass loss rate carbon star \\object{CW Leo} (or \\object{IRC+10216}).  This suggests an isotropic mass loss mechanism with a variability on a broad range of, relatively short, time scales $\\sim$40$-$800\\,yr. Moreover, spatially resolved interferometric CO millimetre observations currently have the potential to trace mass loss rate modulations on the order of a factor 2$-$3 down to a time scale of several hundred years \\citep{Bieging01}. For changes of the mass loss rate over longer periods of time, $\\gtrsim$10$^{5}$\\,yr, statistical studies are generally required.  The emerging picture is that of a mass loss rate which increases gradually with time during the AGB evolution, and its maximum attainable value depends on the initial main-sequence mass \\citep{Habing96}.  Observational results like these are crucial for the understanding of mass loss during late stellar evolution, and for the possibility to pinpoint the mechanism(s) behind this phenomenon, which determines the time scale during the final evolutionary AGB-stage.  \\begin{table*} \\caption{Stellar and circumstellar properties of the sample stars.} \\label{input} \\begin{center} \\begin{tabular}{lllcccccc} \\hline \\noalign{\\smallskip} & & & \\multicolumn{1}{c}{$D$}    & \\multicolumn{1}{c}{$L_*$}    & \\multicolumn{1}{c}{{$\\dot{M}$}$^{\\mathrm d}$}    & \\multicolumn{1}{c}{$v_{\\mathrm{e}}$} & \\multicolumn{1}{c}{$r_{\\mathrm{p}}$$^{\\mathrm e}$} & \\multicolumn{1}{c}{$h$$^{\\mathrm f}$} \\\\ \\multicolumn{1}{c}{{Source}$^{\\mathrm a}$} & \\multicolumn{1}{c}{{Alt. name}$^{\\mathrm b}$} & \\multicolumn{1}{c}{{Var. type}$^{\\mathrm c}$} & \\multicolumn{1}{c}{[pc]} & \\multicolumn{1}{c}{[L$_{\\sun}$]} & \\multicolumn{1}{c}{[M$_{\\sun}$yr$^{-1}$]} & \\multicolumn{1}{c}{[km\\,s$^{-1}$]} & \\multicolumn{1}{c}{[cm]} & \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} {\\object{V384 Per}}        & & Mira & 560 &8100  &3.5$\\times$10$^{-6}$ &          15.0 & 1.4$\\times$10$^{17}$ & 0.4\\\\ {\\object{CW Leo}}          & {\\object{IRC+10216}} & Mira & 120 &9600  &1.5$\\times$10$^{-5}$ &\t    14.5 & 3.7$\\times$10$^{17}$ & 1.0\\\\ {\\object{RW LMi}}\t      & {\\object{CIT 6}} & SRa  & 440 &9700  &6.0$\\times$10$^{-6}$ &\t    17.0 & 1.9$\\times$10$^{17}$ & 1.4\\\\ {\\object{Y CVn}}           & & SRb  & 220 &4400  &1.5$\\times$10$^{-7}$ &\\phantom{0}8.5 & 2.9$\\times$10$^{16}$ & 0.2\\\\ {\\object{IRAS 15194-5115}} & & Mira:& 600 &8800  &1.0$\\times$10$^{-5}$ &\t    21.5 & 3.2$\\times$10$^{17}$ & 1.5\\\\ {\\object{V Cyg}}           & & Mira & 370 &6200  &1.2$\\times$10$^{-6}$ &\t    11.5 & 8.5$\\times$10$^{16}$ & 1.2\\\\ {\\object{S Cep}}           & & Mira & 340 &7300  &1.5$\\times$10$^{-6}$ &\t    22.0 & 7.5$\\times$10$^{16}$ & 0.4\\\\ {\\object{AFGL 3068}}       & & Mira:& 820 &7800  &1.5$\\times$10$^{-5}$ &\t    14.0 & 3.8$\\times$10$^{17}$ & 1.5\\\\ {\\object{LP And}}\t      & {\\object{IRC+40540}} & Mira & 630 &9400  &1.5$\\times$10$^{-5}$ &\t    14.0 & 3.8$\\times$10$^{17}$ & 0.7\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\smallskip  \\noindent $^{\\mathrm a}$ Parameters taken from \\citet{Schoeier01} except for \\object{IRAS 15194-5115} and \\object{AFGL 3068} which are presented in \\citet{Ryde99} and Woods et al.\\ (in prep.), respectively.\\\\ \\noindent $^{\\mathrm b}$ Other frequently used name.\\\\ \\noindent $^{\\mathrm c}$ A colon (:) indicates uncertain classification. \\\\ \\noindent $^{\\mathrm d}$ A CO abundance of 1.0$\\times$10$^{-3}$ relative to H$_2$ was assumed in deriving the mass loss rates.\\\\ \\noindent $^{\\mathrm e}$ Estimated CO photodissociation radius.\\\\ \\noindent $^{\\mathrm f}$ Dust-grain heating parameter (see text for details). \\end{table*}  Carbon monoxide, CO, is a good tracer of the molecular gas content, and it has been extensively used to determine the properties of the circumstellar envelopes (CSEs) formed by the mass loss, e.g., recently \\citet{Schoeier01} used CO millimeter line observations and a detailed modelling of the emission to determine mass loss rates for a large sample of optically bright carbon stars.  The Infrared Space Observatory (ISO; \\cite{Kessler96}) opened a new window for observing AGB-stars, allowing studies of molecular gas, in their expanding CSEs, much closer to the central stars than was previously possible using mainly ground based radio telescopes \\citep{Cernicharo96, Ryde99}. As an example \\citet{Ryde99} studied the high mass loss rate carbon star \\object{IRAS\\,15194-5115} in several rotational transitions in the ground vibrational states of $^{12}$CO and $^{13}$CO. In this way different radial regions of the CSE were probed, and constraints on the temporal changes in the wind characteristics, in particular the mass loss rate, were obtained.  In this paper the method used by \\citet{Ryde99} is adopted in a study of a number of carbon stars observed by ISO. The radiative transfer analysis is further extended to include also the stellar light reemitted by the surrounding dust, and to investigate its effect on the excitation of CO.  \\begin{figure} \\centering{ \\includegraphics[width=88mm]{fig1.eps} \\caption{The FIR spectrum of \\object{RW LMi}, as observed by ISO, showing CO rotational lines in emission superimposed on the continuum emitted by the dust present around the star.  All rotational transitions within the ground vibrational state in this wavelength region (from $J$$=$14$\\rightarrow$13 to $J$$=$29$\\rightarrow$28) are indicated. In addition, the interstellar [CII] line is marked.}} \\label{rwlmi_lws} \\end{figure} ",
        "conclusions": "The modelling of CO rotational line emission at millimetre and FIR wavelengths put constraints on the physical properties of a CSE. Since the available data probe the full radial extent of the CO envelope conclusions about temporal changes in the mass loss rate can be drawn. For the high mass loss rate objects we find that the FIR data, which probe the inner regions of the CSE, are consistent with the results obtained from the radio data.  Under the assumptions of a constant mass loss rate the combined set of data better constrain the envelope parameters such as the mass loss rate and the kinetic temperature of the gas.  From the CO data alone, it is generally hard to put good constraints on any modulations of the mass loss rate. However, analysis of the dust emission put further constraints on the mass loss rate history, and allows conclusions about its temporal changes to be drawn. We find that any longer-term mass loss rate modulations are likely to have been smaller than a factor of $\\sim$5 over the past $\\la$10$^4$\\,yr.  The role of dust in the excitation of CO has been investigated and found to be of only minor importance."
    },
    "0206/astro-ph0206287_arXiv.txt": {
        "abstract": "We study the behaviour of cosmic string networks in contracting universes, and discuss some of their possible consequences. We note that there is a fundamental time asymmetry between defect network evolution for an expanding universe and a contracting universe. A string network with negligible loop production and small-scale structure will asymptotically behave during the collapse phase as a radiation fluid. In realistic networks these two effects are important, making this solution only approximate. We derive new scaling solutions describing this effect, and test them against high-resolution numerical simulations. A string network in a contracting universe, together with the gravitational radiation background it has generated, can significantly affect the dynamics of the universe both locally and globally. The network can be an important source of radiation, entropy and inhomogeneity.  We discuss the possible implications of these findings for bouncing and cyclic cosmological models. ",
        "introduction": "Introduction}  Cosmological scenarios involving oscillating or cyclic universes have been know for a long time \\cite{Tolman}, with interest in them varying according to the latest theoretical prejudices or observational constraints.  Recent interest has been associated with a cyclic extension of the ekpyrotic scenario \\cite{Steinhardt}. A related result was the realization \\cite{Kanekar,Peter1,Peter2} that the presence of a scalar field seems to be necessary to make cosmological scenarios with a bounce observationally realistic. And if scalar fields are present, then one should contemplate the possibility of topological defects being formed.  It is thought that the early universe underwent a series of phase transitions, each one spontaneously breaking some symmetry in particle physics and giving rise to topological defects of some kind \\cite{Kibble1,Book}, which in many cases can persist throughout the subsequent evolution of the universe.  In the present work we study cosmic string evolution in a collapsing universe, following up on and generalizing the results of \\cite{Contracting}, and discuss in much greater detail some implications of the presence of cosmic strings (and cosmic defects in general) for bouncing universes. In a bouncing universe scenario the properties of the universe in the expanding phase depend on physics happening in a previous collapsing phase (before the bounce). For this reason, if defects do exist in these models, it is crucial to understand their evolution and consequences in both the expanding and collapsing phases. Up to now all these studies, be they analytic \\cite{Kibble,Austin,Martins1,Thesis,Martins3} or numerical \\cite{Bennett,Allen1,Moore}, have only been undertaken for the expanding case, and it is clear that while some results may be expected to carry over to the contracting phase, some others clearly won't. This will have consequences not only for the standard string seeded (or hybrid) structure formation scenario \\cite{Avelino2}, but also for other `non-standard' scenarios involving defects, such as the production of adiabatic and nearly Gaussian density fluctuations \\cite{Adiabatic,Gaussian}, or those involving anisotropic or inhomogeneous universes \\cite{Fossils,Inhomog}.  In particular, we expect that cosmic strings will become ultra-relativistic, behaving approximately like a radiation fluid. This means that a cosmic string network, both directly and through the gravitational radiation emitted by the small loops it produces, will soon become a significant source of entropy (and also of inhomogeneity). Hence a cosmic string network is a further problem for cyclic universes if a suitable and efficient mechanism for diluting the entropy is not available.  We should point out at the outset that if/when during the collapse phase one reaches the Hagedorn temperature, one expects the string network to quickly dissolve. However this is largely irrelevant for the points being made in this paper: the radiation, entropy and anisotropy produced by the network will obviously still be left behind if the network does dissipate. On the other hand, it need not be the case that the Hagedorn temperature is reached or, more specifically, that the collapse has to continue all the way to the 'Big Crunch' (where there is no known sensible description of the physics involved). Cosmological models do exist where the bounce takes place at finite size. Indeed, such models seem to be relatively common in scenarios with extra dimensions and scalar fields, although this issue is somewhat debatable.  The outline of this paper is as follows. In Sect. \\ref{cyc} we present a very brief overview of previous work on cyclic universes, and discuss further motivation for our work. Then in Sect.\\ref{stringev} we introduce the basic dynamical properties of cosmic string networks, and after a `warm-up' example of a circular string loop we successively discuss analytic scaling solutions in various regimes in a contracting universe. In Sect. \\ref{numerics} we describe high-resolution numerical simulations of string networks in contracting universes using the Allen-Shellard string code \\cite{Allen1}, and then compare and contrast our analytic and numerical results. Possible cosmological consequences of these results are discussed in Sect. \\ref{conseq}, and finally we present some conclusions in Sect. \\ref{conc}. Throughout the paper we will use units in which $c=\\hbar=1$. ",
        "conclusions": ""
    },
    "0206/astro-ph0206002_arXiv.txt": {
        "abstract": "The ``P32'' Astronomical Observation Template (AOT) provided a means to map large areas of sky (up to 45$\\times$45 arcmin) in the FIR at high redundancy and with sampling close to the Nyquist limit using the ISOPHOT C100 (3$\\times$3) and C200 (2$\\times$2) detector arrays on board ISO. However, the transient response behaviour of the Ga:Ge detectors, if uncorrected, can lead to severe systematic photometric errors and distortions of source morphology on maps. Here we describe the basic concepts of an algorithm which can successfully correct for transient response artifacts in P32 observations. Examples are given demonstrating the photometric and imaging performance of ISOPHOT P32 observations corrected using the algorithm. ",
        "introduction": "From the point of view of signal processing and photometry diffraction-limited mapping in the FIR with cryogenic space observatories equipped with photoconductor detectors poses a particular challenge. In this wavelength regime the number of pixels in detector arrays is limited in comparison with that in mid- and near-IR detectors. This means that more repointings needed to map structures spanning a given number of resolution elements. Due to the logistical constraints imposed by the limited operational lifetime of a cryogenic mission, this inevitably leads to the problem that the time scale for modulation of illumination on the detector pixels becomes smaller than the characteristic transient response timescale of the detectors to steps in illumination. The latter timescale can reach minutes.  Unless corrected for, the transient response behaviour of the detectors will lead to distortions in images, as well as to systematic errors in the photometry of discrete sources appearing on the maps. In general, these artifacts become more severe and more difficult to correct for at fainter levels of illumination, since the transient response timescales increase with decreasing illumination. Compared to the IRAS detectors, the ISOPHOT-C detectors (Lemke et al. 1996) on board the Infrared Space Observatory (ISO; Kessler et al. 1996) had relatively small pixels designed to provide near diffraction limiting imaging. ISOPHOT thus generally encountered larger contrasts in illumination between source and background than IRAS did, making the artifacts from the transient response more prominent, particularly for fields with faint backgrounds.  A further difficulty specific to mapping in the FIR with ISO was that, unlike IRAS, the satellite had no possibility to cover a target field in a controlled raster slew mode. This limited the field size that could be mapped using the spacecraft raster pointing mode alone, since the minimum time interval between the satellite fine pointings used in this mode was around 8\\,s. This often greatly exceeded the nominal exposure time needed to reach a required level of sensitivity (or even for many fields the confusion limit). Furthermore, the angular sampling and redundancy achievable using the fine pointing mode in the available time was often quite limited, so that compromises sometimes had to be made to adequately extend the map onto the background.  A specific operational mode for ISO - the ``P32'' Astronomical Observation Template (AOT) - was developed for the ISOPHOT instrument to alleviate these effects (Heinrichsen et al. 1997). This mode employed a combination of standard spacecraft repointings and rapid oversampled scans using the focal plane chopper. The technique could achieve a Nyquist sampling on map areas of sky ranging up to 45$\\times$45 arcmin in extent (ca. 70$\\times$70 FWHM resolution elements) on timescales of no more than a few hours. In addition to mapping large sources, the P32 AOT was extensively used to observe very faint compact sources where the improved sky sampling and redundancy alleviated the effects of confusion and glitching.  In all, over 6$\\%$ of the observing time of ISO was devoted to P32 observations during the 1995-1998 mission, but the mode could not until now be fully exploited scientifically due to the lack of a means of correcting for the complex non-linear response behaviour of the Ge:Ga detectors. Here we describe the basic concept of a new algorithm which can successfully correct for the transient response artifacts in P32 observations. This algorithm forms the kernel of the ``P32TOOLS'' package, which is now publically available as part of the ISOPHOT Interactive Analysis package PIA (Gabriel et. al 1997; Gabriel \\& Acosta-Pulido 1999). Information on the algorithm, as well as the first scientific applications, can also be found in Tuffs et al. (2002). The user interface of P32TOOLS is described by Lu et al. (this volume). After a brief overview of relevant aspects of the P32 AOT in Sect.~2, we describe the semi-empirical model used to reproduce the transient response behaviour of the PHT-C detectors in Sect.~3. Sect.~4 describes the algorithms used to correct data. Examples demonstrating the photometric and imaging performance of ISOPHOT P32 observations are given in Sect.~5, based on maps corrected using P32TOOLS. ",
        "conclusions": ""
    },
    "0206/astro-ph0206234_arXiv.txt": {
        "abstract": "Dynamic velocity dispersion and mass estimates are given for a sample of five X-ray luminous rich clusters of galaxies at intermediate redshifts ($z\\sim 0.3$) drawn from a sample of 39 clusters for which we have obtained gravitational lens mass estimates. The velocity dispersions are determined from between 9 and 20 redshifts measured with the LDSS spectrograph of the William Herschel Telescope, and virial radii are determined from imaging using the UH8K mosaic CCD camera on the University of Hawaii 2.24m telescope.  Including clusters with velocity dispersions taken from the literature, we have velocity dispersion estimates for 12 clusters in our gravitational lensing sample. For this sample we compare the dynamical velocity dispersion estimates with our estimates of the velocity dispersions made from gravitational lensing by fitting a singular isothermal sphere profile to the observed tangential weak lensing distortion as a function of radius. In all but two clusters, we find a good agreement between the velocity dispersion estimates based on spectroscopy and on weak lensing. ",
        "introduction": "Rich clusters of galaxies are probably the largest well-defined and gravitationally bound objects in the universe, and knowledge of their properties is not only interesting in itself, but can also set important constrains on models of large-scale structure formation. In particular, the cluster mass function has caught major interest (e.g., Bahcall \\& Cen 1993; Gross et al.\\ 1998). Reliable estimates of cluster masses are also important in order to constrain the ratio of baryonic to total mass and to determine the density parameter $\\Omega_0$ (e.g., White \\& Frenk 1991; Lilje 1992; White, Efstathiou, \\& Frenk 1993; Carlberg et al.\\ 1996; Eke, Cole, \\& Frenk 1996).  Cluster masses have been measured using different methods, having different biases and systematics. The method of cluster mass measurement with the longest history is the application of the virial theorem to positions and velocities of cluster galaxies. The first estimates of cluster masses were obtained by \\citet{Zwicky} and \\citet{Smith}. They found that the total mass of galaxies in a cluster only accounted for a small portion of the total cluster mass. This method works well for low redshift clusters, where it can be based on a large number of galaxies, but has recently been applied also to galaxies at intermediate redshifts (e.g., Carlberg et al.\\ 1996) where the number of galaxies observed per cluster is usually relatively small.  Other methods for mass measurement are based on analysis of gravitational lensing by the cluster and consequent distortion of background galaxies and measurements of the X-ray temperature of the hot intra-cluster medium. Of these, weak gravitational lensing has recently become  frequently used for estimating the masses of clusters at intermediate and high redshifts (see e.g., recent reviews by Mellier [1999] and Bartelmann \\& Schneider [2001]). To understand systematic differences between the different mass estimators, comparative studies are necessary. Dahle et al.\\ (2002; hereafter Paper I) presents weak lensing measurements of 39 highly X-ray luminous galaxies at redshift $0.15<z<0.35$ where the masses and their distributions have been determined homogeneously by up-to-date methods. In this paper we present velocity dispersion and virial mass estimates of five clusters from the sample of Paper I, which are compared with the weak lensing results. Our mass determinations are based on the method pioneered by \\citet{LimberMathews} where the mass estimator is \\begin{equation} M_V= \\frac{3\\pi}{2}\\frac{\\sigma_P^2R_{H}}{G}. \\label{eq:vmass} \\end{equation} Here  $\\sigma_P$ is the one-dimensional velocity dispersion of the cluster, while $R_H$ is the projected mean harmonic point-wise separation (projected virial radius) defined by \\begin{equation} R_H=\\frac{N(N-1)}{\\sum_{i<j}R_{ij}^{-1}}, \\label{eq:Rv} \\end{equation} where $N$ is the number of galaxies in the cluster and $R_{ij}$ is the projected distance between any pair of galaxies. In our analysis we have used this method with the modifications of \\citet{Carlberg96} which are described later in the paper. To be able to use equation~(\\ref{eq:vmass}) or versions of it for estimating masses, one has to both determine the velocity dispersion of a cluster through spectroscopy and the angular distribution of cluster galaxies from photometric imaging.   In \\S 2 we present our cluster sample and describe our observing and data reduction procedures for the spectroscopic and photometric observations of the sample.  In \\S 3 we determine the velocity dispersions of the clusters, while we in \\S 4 use the photometric data to estimate the virial radii of the clusters. The virial mass estimates are presented in \\S 5. Finally, in \\S 6, we add to our sample a sample of seven clusters for which we have obtained weak lensing data and for which there exists velocity dispersions in the literature. The velocity dispersions of the clusters of this larger sample are compared with our weak lensing results. ",
        "conclusions": "\\label{chap:discussion}  Our in general good agreement between $\\sigma_P$ determined from dynamical estimates and $\\sigma_{DM}$ determined from gravitational lensing estimates contrast other results.  Smail et al.\\ (1997) presented a comparison of weak lensing and spectroscopic mass measurements for a sample of 9 clusters at $0.17 < z < 0.55$ which were imaged with the WFPC2 instrument aboard the Hubble Space Telescope (HST).  They compared the average tangential shear observed within an annulus around the optical cluster center ($60 < r < 200 h^{-1}\\,{\\rm kpc}$) to the expected shear from SIS-type mass profile with the observed galaxy velocity dispersion. The shear signal was found to be significantly lower than predicted from the spectroscopic measurements for all the clusters, and Smail et al.\\ (1997) suggested that this was caused by systematic overestimates (by typically $\\sim 50$\\%) of the cluster velocity dispersion values given in their paper. As noted in \\S\\ref{Chap:comp}, we see no such discrepancy in our data.  Hoekstra et al.\\ (2001) argue that this discrepancy is instead a result of the small field size of the single WFPC2 pointings used by Smail et al.\\ (1997) to measure the gravitational shear. If there is significant substructure in the central regions of the clusters, or if the density profiles of the clusters are significantly shallower than a SIS-type model at small radii, the shear measurements are likely to underestimate the true cluster mass.  Brainerd et al.\\ (1999) have used ray-tracing through high-resolution $N$-body simulations of massive ($M_V \\sim 10^{15}h^{-1}M_{\\sun}$) clusters to show that the average tangential shear, when measured within an annulus with an outer radius $R_{\\rm MAX}$ significantly smaller than the virial radius (and assuming a SIS-type mass profile), will tend to underestimate the true cluster mass. For $R_{\\rm MAX} \\simeq R_V$, they find that the SIS-model fit yields quite accurate mass values (typically 5\\% -- 10\\% underestimates; see their Figure~9) even if the true mass density profile of the clusters is more accurately represented by the model proposed by Navarro, Frenk, \\& White (1995). We note that for most of the clusters listed in Table~\\ref{tab:cluster_data}, we have estimated $\\sigma_{\\rm DM}$ based on tangential shear measurements out to radii $\\simeq R_V$ (for most of our clusters, we have used $R_{\\rm MAX} = 550\\arcsec = 1.5 h^{-1}\\,{\\rm Mpc}$ at $z = 0.3$; see Paper I for details), so we expect our values to be essentially unbiased.  Girardi \\& Mezzetti (2001) have reanalyzed the available radial velocity measurements in the clusters studied by Smail et al.\\ (1997) and derive velocity dispersion values that are on average about 20\\% lower than those given by Smail et al.\\ (1997). Thus, it appears that the discrepancy between spectroscopic and lensing mass estimates seen by Smail et al.\\ (1997) is most probably caused by a simultaneous overestimation of dynamical masses and underestimation of lensing masses, of almost equal magnitude.  Girardi \\& Mezzetti (2001) also compare their virial cluster mass estimates with weak lensing mass estimates drawn from the literature for 18 clusters (for 7 of these the lensing mass was measured from single HST WFPC2 pointings), and find a median value $M_{spec}/M_{lens} = 1.30$ and a 90\\% confidence interval $0.63 < M_{spec}/M_{lens} < 2.13$. This is consistent with our results in \\S\\ref{Chap:comp} within the errors, but the $M_{lens}$ values they use are probably subject to a larger bias caused by the (on average) significantly smaller fields used to measured the shear.  In their study, Brainerd et al.\\ (1999) did not consider the possible effect of structures outside their clusters on the lensing mass estimates.  Galaxy clusters are believed to be connected by large filaments of dark matter, diffuse gas, and galaxies, and mass within such filaments will contribute significantly to the lensing mass when seen in projection along the line of sight to a cluster, even if it is physically separated from the cluster by several Mpc (Metzler, White, \\& Loken 2001).  Still, even in this scenario, $\\sigma_{\\rm DM}$ may not necessarily be significantly larger than $\\sigma_P$, since the galaxy velocity dispersion would also be biased upwards by galaxies outside the virial radius falling into the main cluster along the filaments.  The modest fraction ($\\sim 30$\\%) of massive clusters that have significant substructure in their dark matter distribution provide evidence that the majority of massive clusters at $z = 0.2$--0.3 are close to dynamical equilibrium (Dahle et al.\\ 2002a).  The good agreement that we find between virial masses and lensing masses for most of our clusters is consistent with this picture, although significant amounts of substructure could probably be present in the clusters without changing the $\\sigma_P$ values beyond the error bounds of our measurements. To reliably assess the dynamical state of the clusters from spectroscopic data alone, it would probably be necessary to observe galaxy samples an order of magnitude larger than those we present in this paper. In addition, X-ray observations, particularly using Chandra and XMM, would be an efficient aid to measure substructure and determine the dynamical state of the clusters.  We plan to investigate the relation between weak lensing and X-ray properties for a larger sample of clusters in a future paper (H.\\ Dahle et al.\\ 2002b, in preparation)."
    },
    "0206/astro-ph0206372_arXiv.txt": {
        "abstract": "We analyze the evolution of energy density fluctuations in cosmological scenarios with a mixture of cold dark matter and quintessence, in which the quintessence field is modeled by a constant equation of state. We obtain analytic expressions for the time evolution of the quintessence perturbations in models with light fields.  The fluctuations behave analogously to a driven harmonic oscillator, where the driving term arises from the inhomogeneities in the surrounding cosmological fluid.  We demonstrate that the homogeneous solution, determined by the initial conditions, is completely sub-dominant to the inhomogeneous solution for physically realistic scenarios. Thus we show that the cosmic microwave background anisotropy predicted for such models is highly insensitive to the initial conditions in the quintessence field. ",
        "introduction": "\\label{secintro} \\mynote{WHY Q,WHATS Q} Growing observational evidence suggests that the total matter density of the universe is significantly less than the critical density. Yet, measurements of the cosmic microwave background (CMB) indicate that the universe is flat and the total energy density is precisely equal to the critical density \\cite{Ostriker}. Candidates for the missing energy are the cosmological constant ($\\Lambda$) \\cite{LambdaRef}, and quintessence ($Q$), a time-varying, spatially inhomogeneous component with negative pressure \\cite{PJS,Cald98}. Examples of quintessence include slowly evolving fields or topological defects such as a network of light and tangled cosmic strings \\cite{Spergel}-\\cite{kess}. Cold Dark Matter (CDM) models with quintessence (\\qcdm) which fit the data from observations of high red shift supernovas, gravitational lensing, CMB anisotropy,  and structure  formation \\cite{Ostriker}-\\cite{PJS} have been found.  \\mynote{WHICH MODELS,WHAT'S INSENSITIVITY?} For the purposes of this paper, we model quintessence as a scalar field quintessence rolling down a potential $V(Q)$ with an equation of state $w \\equiv p_Q / \\rho_Q$, where $p_Q$ is the pressure and $\\rho_Q$ is the energy density of the field. We consider here a large class of models in which the field and fluctuation evolution may be represented by a constant equation of state with value between 0 and -1, and in which the sound speed in quintessence fluctuations, $c_{sQ}^2$, approaches unity at scales much smaller than the horizon. Most models of quintessence that have appeared in the literature satisfy these conditions. This class includes not only ``tracker models'', which have dynamical attractor behavior, but also more general potentials \\cite{Weiss}-\\cite{zlatev}. We show that, in such models, the CMB anisotropy is insensitive to initial conditions of the quintessence field. By insensitive we mean that the fractional change in the CMB anisotropy power spectrum due to a change in initial conditions cannot be observationally resolved.  These conclusions were originally mentioned in our first paper on quintessence models \\cite{Cald98}, and also by others \\cite{Viana:1997mt, Perrottaetal, Braxetal}, who noted that the amplitude of the energy density perturbations in the background fluid are largely independent of the initial conditions in the scalar field, provided the initial energy contrast is less than unity. Our present work is consistent with these earlier conclusions. However, we go further in this paper by examining in detail the behavior of the scalar field perturbations, and the causes for the insensitivity, as described next.  \\mynote{ROADMAP}In the next section we derive the equations of motion for the evolution of the \\qfield and its fluctuations. We introduce a formalism to numerically study the % quintessence fluctuations in terms of the evolution of the equation of state as a function of cosmological scale-factor. In section \\ref{section.approx} we show that models with constant equation of state represent the behavior of a large class of models with light fields and monotonically evolving equations of state. In section \\ref{section.evol} we analytically solve the fluctuation equation for constant equation of state at large wavelengths. We use this solution to describe key features of the fluctuation evolution obtained from the numerical integration. Finally, in section \\ref{secinits} we report the effect of changing initial conditions on quintessence fluctuation evolution and CMB anisotropy. We show that the anisotropy is highly insensitive to changes in the initial conditions. \\mynote{For example, consider models with $w > -0.9$ and long wavelength matter fluctuations with amplitude $\\delta \\rho_m / \\rho_m \\sim 10^{-5}$ at horizon re-entry (which corresponds to an an initial value of $\\delta \\rho_m / \\rho_m \\sim 10^{-16}$ at $a=10^{-8}$ in the synchronous gauge). Then, for any initial value of $\\delta \\rho_Q/\\rho_Q$ less than $10^{11}$ in the synchronous gauge, there is no distinguishable difference in the CMB anisotropy.} ",
        "conclusions": "\\label{sec.concl}  We studied in this paper a large class of quintessence models with light fields and sound speed $c_s^2 \\sim 1$ at small wavelengths, which have the property that they can be well approximated by constant equation of state, $w$. The evolution of the fluctuations in these models was obtained by numerical integration and explained by approximate analytic solutions to the fluctuation equation at large wavelengths. Our central result is that the CMB anisotropy in such models is insensitive to initial conditions on the quintessence fluctuations for smooth and adiabatic initial conditions . For $w=-0.9$, the CMB anisotropy is insensitive in the large range of initial conditions $\\Big(\\qfrac\\Big)_{init} < 10^{11} \\,(F=10^5)$ for $\\Big(\\mfrac\\Big) \\sim 10^{-5}$ at horizon re-entry. Secondly, the sensitivity increases as $w$ approaches -1. At $w=-0.999$, the range reduces to $\\Big(\\qfrac\\Big)_{init} < 10^{9} \\,(F=10^3)$. However, physically reasonable models such as those based on inflation and ekpyrosis do not produce such large values of $\\qfrac$, and the ratio of energy in quintessence fluctuations to that in matter fluctuations is much smaller than unity. Hence, we do not anticipate that the CMB anisotropy will be sensitive to initial conditions in realistic cases. The same analytical arguments made in this paper carry over to the more general quintessence models in which $w$ is more strongly time dependent or $c_s^2 \\ne 1$ at small wavelengths. However, the precise numerical lower bound on the initial conditions required to imprint a distinguishable effect on the CMB anisotropy has to be worked out on a case by case basis.  \\medskip  This work was supported by DOE grant DE-FG02-95ER40893 (RD), NSF grant PHY-0099543 (RRC), and DOE grant DE-FG02-91ER40671 (PJS).  \\noindent"
    },
    "0206/astro-ph0206329_arXiv.txt": {
        "abstract": "{ We have carried out an extensive search for X-ray emission from young, very low-mass objects near and beyond the substellar limit, making use of archived {\\em ROSAT} PSPC and HRI observations pointed at Brown Dwarfs and Brown Dwarf candidates in the young $\\sigma$\\,Orionis and Taurus-Auriga associations. In \\sO \\ we identify three Brown Dwarf candidates with X-ray sources; in Taurus-Auriga we add one further X-ray detection of a Brown Dwarf to the list published earlier. We combine this data with all previously X-ray detected Brown Dwarfs and Brown Dwarf candidates in young stellar associations and star forming regions to perform a study of stellar activity parameters on the as yet largest sample of young, very low mass objects. A similar relation between X-ray and bolometric luminosity, and H${\\alpha}$ emission, respectively, as is known for T Tauri stars seems to hold for young objects down to the substellar limit, too. No signs for a change in X-ray activity are found on the transition to substellar masses. ",
        "introduction": "\\label{sect:intro}  Late-type stars exhibit strong signs of magnetic activity such as H$\\alpha$, Ca\\,II or X-ray emission from hot thermal plasma confined in magnetic fields on the star. For fully convective stars (with spectral type $\\sim$\\,M3 and later) the change in interior structure is expected to result in a change of the field sustaining dynamo and, therefore, of the emission properties. However, no clear change in common activity diagnostics is found at these spectral types. \\citey{Gizis00.1}, \\citey{Basri01.1} and \\citey{Mohanty02.1} have observed a decline in H$\\alpha$ emission for old ($>1$\\,Gyr) L-type objects in the field, suggesting that a decline in dynamo activity sets in beyond the boundary where the objects become fully convective and close to the borderline to substellar masses. An X-ray study of K- and early M-stars within the solar neighborhood by \\citey{Fleming95.1} did not unveil any change in X-ray activity for objects on the transition towards fully convective energy transport. It is unclear to date how activity in fully convective and substellar objects depends on parameters such as rotation or age. Studying these relations is essential to an understanding of the underlying dynamo mechanism.  An important measure for magnetic activity on late-type stars of all ages is X-ray emission which is usually explained by emission from a hot plasma. Within the last few years, X-ray emission has also been detected from Brown Dwarfs (BD) and BD Candidates (see e.g. Neuh\\\"auser \\& Comer\\'on 1998, Neuh\\\"auser et al 1999). In these early studies with {\\em ROSAT} only BDs in the Chamaeleon, Taurus and $\\rho$\\,Ophiuchus star forming regions were detected, while all older substellar objects in stellar associations such as the Pleiades and in the field are X-ray quiet down to the {\\em ROSAT} detection limit. More recent observations with {\\em Chandra} revealed X-ray emission from further young BDs in Orion (\\cite{Garmire00.1}, \\cite{Feigelson02.1}), $\\rho$\\,Oph (\\cite{Imanishi01.1}), and IC\\,348 (\\cite{Preibisch01.1}; \\cite{Preibisch02.1}). To date only one old field BD is known to emit X-rays, namely LP\\,944-20 which was detected only during a flare (\\cite{Rutledge00.1}). Here, we extend the investigations with a study of X-ray emission from the very low-mass (VLM) members in the $\\sigma$\\,Orionis association, and an update of the X-ray emitting BD population in Taurus.  The $\\sigma$\\,Ori star forming region was discovered by \\citey{Walter97.1}, and was shown to be rich in X-ray sources. Photometric and spectroscopic observations by \\citey{Bejar99.1}, \\citey{Zapatero00.1}, and \\citey{Bejar01.1} have revealed $\\sim$\\,80 objects in the low mass regime close to and below the hydrogen burning mass limit (HBML) in the OB1b association near the multiple star $\\sigma$\\,Orionis at a mean distance modulus of $DM = 7.73$ measured by {\\em HIPPARCOS}, corresponding to $352$\\,pc (\\cite{Bejar99.1}). The age of this cluster is estimated to be $1-7$\\,Myrs where the upper age limit is given by constraints of the central star $\\sigma$\\,Orionis: being of spectral type O9.5 and still in the hydrogen burning phase, it cannot be older than $5-7$\\,Myrs (see \\cite{Bejar01.1}, \\cite{Barrado01.1}, and references therein). According to \\citey{Bejar01.1}, the position of all members in the H-R diagram is best compatible with an isochrone for 5\\,Myrs. Based on measurements of lithium abundances in \\sO \\ members \\citey{Zapatero02.1} give as most likely age $2-4$\\,Myrs with an upper limit of 8\\,Myrs. At this young age, even VLM objects are still very luminous. In addition, the $\\sigma$\\,Ori association is hosted in a region of very low extinction ($E_{B-V}=0.05$; \\cite{Lee68.1}). This provides excellent conditions to study activity in faint VLM objects despite the considerable distance to the cluster.  Taurus-Auriga is one of the nearest ($d$ = 140\\,pc; \\cite{Elias78.1}) and most-studied regions of star formation. The X-ray emission from late-type pre-main sequence stars in Taurus-Auriga was recently discussed by \\citey{Stelzer01.1} (hereafter SN01). SN01, whilst concentrating on the G- to early M-type T Tauri Stars, put forth the {\\em ROSAT} PSPC detection of several objects with spectral types beyond M5, i.e. near the substellar limit. Four additional BDs have been discovered in that region (\\cite{Martin01.2}) since. In this paper we present a detailed analysis of the X-ray activity of all VLM objects in Taurus-Auriga including the new BDs. In the following we used the term 'very-low mass' for objects with spectral type later than M4.  In Sect.~\\ref{sect:sample}, we outline the criteria we used when selecting our data samples. In Sect.~\\ref{sect:data}, we describe the {\\em ROSAT} data analysis. We present the results for both $\\sigma$\\,Orionis and Taurus in Sect.~\\ref{sect:results}. In Sect.~\\ref{sect:discussion} we perform a comparative study of the X-ray properties of all young BDs and BD candidates detected so far, including tests for variability and an investigation of correlations with other activity parameters. A summary of our results is given in Sect.~\\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}  We have searched for X-ray emission from VLM objects near and below the substellar limit in the young stellar associations \\sO \\ and Tau. Three objects in \\sO \\ and 13 in Tau were detected in {\\em ROSAT} PSPC and HRI observations. We combined these results with all X-ray detections from BDs and BD candidates in star forming regions available up to date and studied stellar activity parameters such as X-ray emission, the ratio $\\lg {\\frac{L_{\\rm x}}{L_{\\rm bol}}}$, H$\\alpha$ emission and X-ray hardness ratios of the whole sample. The comparison of the X-ray emission of the young BDs and BD candidates with that of (higher-mass) TTS in Taurus shows that $L_{\\rm x}$ decreases monotonically into the BD regime. However, $\\lg {\\frac{L_{\\rm x}}{L_{\\rm bol}}}$ remains approximately constant suggesting that the efficiency at which hot coronal gas is produced does not change. For a statistical evaluation of X-ray variability we applied the KS-test to all {\\em ROSAT} detected BDs and BD candidates in star forming regions. We found that $\\sim 17$\\,\\% of the sources are variable with $> 95$\\,\\% confidence. To complete the study on activity parameters in young VLM stellar and substellar objects, there is a clear need for rotation and  H$\\alpha$ measurements and a search for infrared excess giving clues for the presence of circumstellar disks."
    },
    "0206/astro-ph0206435_arXiv.txt": {
        "abstract": "We present spectroscopy and time-series photometry of the dwarf nova QZ Ser.  The spectrum shows a rich absorption line spectrum of type $K4 \\pm 2$.  K-type secondary stars are generally seen in dwarf novae with orbital periods $P_{\\rm orb} \\sim 6$ h, but in QZ Ser the absorption radial velocities show an obvious modulation (semi-amplitude $207(5)$ km s$^{-1}$) at $P_{\\rm orb} = 119.752(2)$ min, much shorter than typical for such a relatively warm and prominent secondary spectrum. The H$\\alpha$ emission-line velocity is modulated at the same period and roughly opposite phase. Time-series photometry shows flickering superposed on a modulation with two humps per orbit, consistent with ellipsoidal variation of the secondary's light. QZ Ser is a second example of a relatively short-period dwarf nova with a surprisingly warm secondary.  Model calculations suggest that the secondary is strongly enhanced in helium, and had already undergone significant nuclear evolution when mass transfer began.  Several sodium absorption features in the secondary spectrum are unusually strong, which may indicate that the present-day surface was the site of CNO-cycle hydrogen burning in the past. ",
        "introduction": "Cataclysmic variable stars (CVs) are close binary systems in which a low-mass secondary transfers mass onto a white dwarf; \\citet{warn} wrote an excellent monograph on CVs.  The Roche geometry tightly constrains the secondary star's mass at a given orbital period $P_{\\rm orb}$.  Short-period systems have low-mass secondaries, so if the chemical composition is normal ($X \\sim 0.7$), the secondary is a faint M dwarf or brown dwarf and contributes negligibly to the visible-light spectrum (Fig.~4 of \\citealt{patprecess01}); However, \\citet{thor02} found a K4 $\\pm 2$ secondary in the dwarf nova 1RXS J232953.9+062814 (hereafter RX 2329+06), which has $P_{\\rm orb}$ = 64 min. They suggested that the secondary was somewhat evolved at the start of mass transfer, with its core substantially enhanced in helium.  In this scenario the portion of the secondary remaining today corresponds to the core of the original star, and the enhanced helium greatly affects the mass-temperature relation.  Here we present new observations of the dwarf nova QZ Ser, which appears to be a close relative of RX 2329+06. QZ Ser was discovered by Katsumi Haseda in 1998 and designated as HadV04 in the discovery notification (vsnet-obs 18349)\\footnote {The vsnet mailing list archives are available at http://vsnet.kusastro.kyoto-u.ac.jp/vsnet/index.html} , in which T. Kato suggested an identification with the ROSAT WGA source 1WGAJ1556.9+2108 \\citep{wga94}.  The outbursting object was discovered on small-scale patrol films, and its quiescent counterpart was initially uncertain.  On 2001 June 28 we obtained a spectrum of the star nearest the position, and it proved to be an ordinary late-type star. Returning to the field on 2002 January 20 we obtained a spectrum of a somewhat fainter star 11 arcsec to the northwest. This showed the distinctive broad Balmer emission characteristic of dwarf novae at minimum light, together with absorption features of a late-type star.   On 2002 Feb 3.23 UT, P. Schmeer detected an outburst of QZ Ser (vsnet-campaign 1281), and an examination of an outburst image by H. Yamaoka (vsnet-alert 7172) confirmed that the outbursting object was on the position of the emission-line object, cementing the identification.  The position of QZ Ser, derived from a fit to 62 USNO A2.0 stars \\citep{mon96} on one of our 1.3m images (described below), is $\\alpha_{\\rm ICRS} = 15^{\\rm h}\\ 56^{\\rm m}\\ 54^{\\rm s}.50,\\ \\delta_{\\rm ICRS} = +21^{\\circ}\\ 07'\\ 19''.5\\ (\\pm 0.''3$ estimated uncertainty).  Its position in the USNO A2.0 is not significantly different, which sets an upper limit on its proper motion $\\mu \\simle 0''.03$ yr$^{-1}$. The {\\it Living Edition} of the {\\it Catalog and Atlas of Cataclysmic Variables} \\citep{downes2001} gives an up-to-date finding chart. ",
        "conclusions": "QZ Ser is only the second short-period dwarf nova to show a K-star secondary. In all other dwarf novae with $P_{\\rm orb} \\simle 2$ hr, save for a handful of helium systems, the secondaries are late M dwarfs which contribute a small fraction of the visible light.  How are we to account for this unusual object? As with RX 2329+06 \\citep{thor02}, we suggest that the secondary evolved significantly on the main sequence prior to mass transfer, enhancing the core with enough helium to greatly affect the mass-$\\te$ relation.  The small, hot secondary we see is the remnant of a once more massive secondary.  For illustration, we computed some evolutionary models in the framework of the standard disrupted magnetic braking scenario (see \\citealt{bk00} and references therein). At the $\\sim 2$ hr orbital period of QZ Ser, a donor which started mass transfer on the zero-age main sequence (ZAMS) would have spectral type M4-M5 and mass $M_2 \\, \\sim \\, 0.2 \\, \\msol$. Such a ``standard\" sequence, based on the models of \\citet{bk00}, is displayed in Fig.~\\ref{fig4} by the solid curve. Also shown in Fig.~\\ref{fig4} are some test calculations in which mass transfer begins near the end of central H burning.   The tracks shown represent several choices of {\\it initial} secondary masses $M_2 \\, \\simgr \\, 1 \\, \\msol$, constant mass loss rates ${\\dot M} \\sim 10^{-9} - 10^{-8} \\, \\msolyr$, and different levels of nuclear evolution $X_{\\rm c} \\simle 0.1$. As Fig.~\\ref{fig4} shows, these naturally reproduce the observed properties of QZ Ser and RX 2329+06. The choice of our parameters seems reasonable and we do not require particularly extreme assumptions on our evolutionary scenarios to fit these objects\\footnote{With our assumptions, the system in principle should have initially tranferred mass on a thermal timescale, until $M_2$ became small enough for the system to reappear as a standard CV (see \\citealt{bk00} for a discussion).}.  As already mentioned in \\citet{thor02}, the models predict altered surface abundances. Fig. \\ref{fig4} displays the surface enrichment of $^{14}$N processed by the CNO cycle in the deeper layers and mixed up to the surface. The surface helium abundance increases also, reaching mass fraction $\\sim$ 0.6 at $ P_{\\rm orb} \\sim 2$h.  In RX 2329+06 the strengths of H$\\alpha$ and HeI $\\lambda$5876 are in the ratio 3.6:1, whereas this ratio is typically 6 and above in SU UMa stars \\citep{thor02}; in the 2002 February spectra of QZ Ser, the ratio is 5:1, which may indicate some enhancement of He.  As noted earlier, though, He features were nearly absent in the 2002 January pre-outburst spectra, so the He:H line ratio is evidently not a consistent measure of abundance in this object.  The sodium enhancement noted earlier (if real) may provide an important window on the secondary's nuclear processing history. The temperatures reached in the deepest layers of stars with masses $\\simgr \\, 1.2 \\, \\msol$ near the end of central H burning are high enough to allow the production of $^{23}$Na via the $^{22}$Ne($p, \\gamma)^{23}$Na reaction.  During mass transfer, the convective envelope proceeds inward and may reach the Na-enriched layers.  At the 2-hour period of QZ Ser, the bottom of the secondary's convective envelope should reach into layers which have attained temperatures of 1.5 -- $2 \\times 10^7$ K during prior evolution. Based on the recent NACRE reaction rates \\citep{angulo99}, this is hot enough to destroy $^{22}$Ne by proton capture, but it is not hot enough for the Na-Ne cycle to contribute significant Na enhancement \\citep{weiss00}. A combination of deep mixing and Ne-Na cycling has been considered as an explanation for anomalous Na abundance in globular cluster giants \\citep{denisenko90, kraft97, weiss00}. The nuclear reaction networks presently implemented in our models unfortunately do not produce a quantitative estimate the sodium enhancement, but work is in progress to remedy this. At present, we can only point to the strong Na lines as a clue that the surface material was processed at relatively high temperatures, mixed upward, and exposed as mass transfer stripped away the overlying layers.  The origin of the Na enhancents in the globular cluster stars is uncertain, and as in that case, it is possible that QZ Ser's sodium was already present in the gas from which the system formed. It would therefore be desirable to detect the primary products of CNO processing -- He and N enrichment, and C and O depletion.  Unfortunately, test model atmospheres with $\\te$ = 4300 K and $\\log g = 5.5$, kindly computed by F. Allard (private communication) showed essentially no observable effect when He was enhanced to 50\\% in mass fraction and the CNO abundances were altered as expected.  More noticeable effects should appear in models at cooler temperatures ($\\te < 4000$ K), where molecules involving C or O form.  If our interpretation based on nuclearly evolved donors is correct, we may expect that a non-negligible number of short period CVs with anomalously hot secondaries will be discovered.  As already emphasized by \\citet{beuermann98} and \\citet{bk00}, such evolved sequences can indeed explain a  {\\it substantial fraction} of the observed CVs with late spectral types and orbital periods $P \\, \\simgr \\, 6$ h. As shown in Fig. \\ref{fig4}, the most evolved sequences which provide an explanation for such systems predict as well the existence of early spectral types at shorter periods. The late spectral type systems with $P \\, \\simgr \\, 6$ h represent a non-negligible fraction of systems in the sample of \\citet{beuermann98}. \\citet{bk00} were thus concerned by the fact that such evolved sequences may remain active in the 2-3 h period gap, since such extreme secondaries never become fully convective. In order to prevent populating the period gap and predicting too early spectral types at shorter periods, (\\citealt{bk00}, see their \\S 4) suggested an increase of the mean mass transfer rate during the secular evolution of such evolved donors. However, the  existence of QZ Ser at $P_{\\rm orb}$ = 2.0 h now supports the idea that evolved sequences may remain active in the gap.  Our proposed scenario would be supported if QZ Ser and RX 2329+06 were found to have enhanced CNO-process abundances at their surfaces, and a more quantitative study of the abundances of Na, Al, and other light metals may provide this evidence despite the low $\\te$.   In the models we have computed to date, the secondaries in longer-period systems are significantly less massive than in `standard' systems.  If this tendency proves robust, accurate measures of donor masses in long-period systems could identify systems destined to evolve into stars like RX 2329+06 and QZ Ser in our scenario. Finally, if our interpretation is correct, one should find similar systems {\\it in the period gap}.  {\\it Acknowledgments.} We gratefully acknowledge funding by the NSF (AST 9987334), and we thank the MDM staff for their support. We thank especially France Allard for computing the test model atmospheres.  \\clearpage"
    },
    "0206/astro-ph0206090_arXiv.txt": {
        "abstract": "{ We present feasibility studies to directly image stellar surface features, which are caused by magnetic activity, with the Very Large Telescope Interferometer (VLTI).  We concentrate on late type magnetically active stars, for which the distribution of starspots on the surface has been inferred from photometric and spectroscopic imaging analysis. The study of the surface spot evolution during consecutive rotation cycles will allow first direct measurements (apart from the Sun) of differential rotation which is the central ingredient of magnetic dynamo processes. The VLTI will provide baselines of up to 200 m, and two scientific instruments for interferometric studies at near- and mid-infrared wavelengths. Imaging capabilities will be made possible by closure-phase techniques. We conclude that a realistically modeled cool surface spot can be detected on stars with angular diameters exceeding $\\sim$\\,2\\,mas using the VLTI with the first generation instrument AMBER. The spot parameters can then be derived with reasonable accuracy. We discuss that the lack of knowledge of magnetically active stars of the required angular size, especially in the southern hemisphere, is a current limitation for VLTI observations of these surface features. ",
        "introduction": "Optical interferometers are about to become powerful instruments to image vertical and horizontal temperature profiles and inhomogeneities of stellar surfaces. Direct observations of these structures will be the key for constraining underlying hydrodynamic and magneto-hydrodynamic mechanisms and for our further understanding of related phenomena like predictions of stellar activity and asymmetric or stochastic mass-loss events. The study of the evolution of surface spots which are caused by magnetic activity, during consecutive rotation cycles will allow first direct measurements (apart from the Sun) of differential rotation which is the central ingredient of magnetic dynamo processes.  Many starspots have been discovered by photometric monitoring (e.g. Strassmeier et al. 1999). The distribution of starspots on the surfaces of stars has so far been inferred from photometric and spectroscopic imaging analysis for about 50 stars (''Summary of Doppler images of stars'', www.aip.de/groups/activity/DI/summary). In addition, there is evidence for surface features on some slowly rotating single K giant stars based on Ca\\,II variability (e.g. Choi et al. 1995). The observation of these stars may not be very feasible for Doppler imaging techniques because of their slow rotation. However, they might be good candidates for interferometric imaging, here because of their slow rotation, and their sufficiently large angular diameters (Hatzes et al. 1997). Magnetic activity is usually assumed to be the cause for these surface spots.  Optical interferometry has already proven its ability to derive stellar surface structure parameters beyond diameters.  The limb-darkening effect on stellar disks was confirmed by interferometry for several stars by Hanbury Brown et al. (1974), Haniff et al. (1995), Quirrenbach et al. (1996), Burns et al. (1997), Hajian et al. (1998), Young et al. (2000), and Wittkowski et al. (2001). The latter measurement did not only succeed in discriminating uniform disks and limb-darkened disks, but also in constraining model atmosphere parameters. For the apparently largest supergiants $\\alpha$~Orionis, $\\alpha$~Scorpii and $\\alpha$~Herculi, bright spots have been detected and mapped by direct imaging techniques including optical interferometry and HST imaging (Buscher et al. 1990, Wilson et al. 1992, Gilliland \\& Dupree 1996, Burns et al. 1997, Tuthill et al. 1997, Young et al. 2000). Large-scale photospheric convection (Schwarzschild 1975, Freytag et al. 1997) is a preferred interpretation for the surface features on the these supergiants, and for interferometric observations of asymmetric features on or near stellar disks (see, e.g. Perrin et al. 1999, Weigelt et al. 2000) or circumstellar envelopes (see, e.g. Weigelt et al. 1998, Wittkowski et al. 1998).  The ESO VLTI (Very Large Telescope Interferometer), a European general user interferometer located on Cerro Paranal in northern Chile, is about to become an excellent facility to study stellar surfaces by means of optical interferometry. It will provide a large choice of baselines up to 200\\,m, the use of the 8\\,m Unit Telescopes (maximum baseline 130\\,m), of several 1.8\\,m Auxiliary Telescopes as well as of a near-infrared instrument (AMBER, see Petrov et al. 2000), and a mid-infrared instrument (MIDI, see Leinert et al. 2000). The AMBER instrument will provide imaging capabilities by closure-phase techniques. For a description of the VLTI and the recent achievement of first fringes see Glindemann et al. (2000), Glindemann et al. (2001a), Glindemann et al. (2001b), and references therein. Since then, commissioning activities have been carried out. VLTI commissioning data of scientific relevance are being publicly released ({\\small www.eso.org/projects/vlti/instru/vinci/vinci\\_data\\_sets.html}). Measurements of visibility amplitudes beyond the first minimum of the visibility function, which is mandatory for an analysis of stellar surface features, became already feasible with the commissioning instrument VINCI (see ESO Press Release 23/01; Wittkowski et al., in preparation). The potential to study stellar surfaces with the VLTI has previously been investigated by von der L\\\"uhe et al. (1996), Hatzes (1997) and von der L\\\"uhe (1997). Other existing or planned ground-based optical and near-infrared interferometric imaging facilities (with location and maximum baseline) include the CHARA (Mt. Wilson, USA, 350\\,m), COAST (Cambridge, UK, 100\\,m), IOTA (Mt. Hopkins, USA, 40\\,m), NPOI (Flagstaff, USA, 437\\,m), KECK (M. Kea, USA, 140\\,m), LBT (Mt. Graham, USA, 22\\,m), OHANA (M. Kea, USA, 800\\,m). In addition, free-flying space-based interferometers are being designed. For instance, the ``stellar imager'' (SI), a NASA design for a space-based UV/optical interferometer with a characteristic spatial resolution of 0.1 mas, is especially intended to image dynamo activity of nearby stars (hires.gsfc.nasa.gov/~si).  Here, we investigate to which extent and accuracy surface features with realistically modeled parameters can be analyzed in the near future by VLTI measurements using first generation instruments. Observing restrictions and accuracy limits are taken into account. We concentrate on magnetically active stars and studies which aim at understanding stellar magnetic dynamo processes. ",
        "conclusions": "We have investigated the feasibility to image surface spots on magnetically active stars with the VLT interferometer. We used realistic model assumptions based on a Doppler image of the giant star CM\\,Cam. Observations were considered which make use of the first generation VLTI instrument AMBER. All observing restrictions were taken into account, as for example the limited lengths of the delay lines. We used conservative assumptions for the accuracy and calibration uncertainty of the AMBER instrument. We have discussed that magnetically active single giants are the most promising candidate stars for interferometric imaging. We show by simulations of observations which are performed during only one night, that the surface feature on our model giant can be detected at a high significance level and estimate the accuracy of fitted spot parameters to be reasonably good. The combination of data taken during one night limits the rotational period of the star to about 15-20 days. For faster rotating stars, the rotation period could be modeled as an additional parameter. Our analysis shows as well that the angular diameter of our model star of 2.6\\,mas is at the lower limit for feasible interferometric studies of stellar surfaces using the VLTI with AMBER. We find, on the other hand, that the majority of stars for which the existence of surface features caused by magnetic activity is known or strongly suggested, have angular diameters smaller than 2\\,mas and, in addition, are located in the northern hemisphere. It would be desirable to initiate spectroscopic and photometric variability studies in the southern hemisphere, focusing on apparently large giants, as a preparation and source of complimentary information for VLTI measurements of stellar surface features. The VLTI, other ground based interferometers, and finally large space based UV interferometers promise to give us exciting new insights into magnetic and hydrodynamic activity of stars, with first results to be expected very soon."
    },
    "0206/astro-ph0206059_arXiv.txt": {
        "abstract": "The Jovian dust streams are high-speed bursts of submicron-sized particles traveling in the same direction from a source in the Jovian system.  Since their discovery in 1992, they have been observed by three spacecraft: Ulysses, Galileo and Cassini.  The source of the Jovian dust streams is dust from Io's volcanoes.  The charged and traveling dust stream particles have particular signatures in frequency space and in real space. The frequency-transformed Galileo dust stream measurements show different signatures, varying orbit-to-orbit during Galileo's first 29 orbits around Jupiter. Time-frequency analysis demonstrates that Io is a localized source of charged dust particles. Aspects of the particles' dynamics can be seen in the December~2000 joint Galileo-Cassini dust stream measurements. To match the travel times, the smallest dust particles could have the following range of parameters: radius: 6~nm, density: 1.35--1.75~g/\\pccm, sulfur charging conditions, which produce dust stream speeds: 220$\\backslash$450~\\kms\\ (Galileo$\\backslash$Cassini) and charge potentials: 5.5$\\backslash$6.3~V (Galileo$\\backslash$Cassini). ",
        "introduction": "The Jovian dust streams are high-speed collimated streams of submicron-sized particles traveling in the same direction from a source in the Jovian system.  They were discovered in March~1992 by the cosmic dust detector instrument onboard the Ulysses spacecraft, when the spacecraft was just past its closest approach to Jupiter. Observations of the Jovian dust stream phenomena continued in the next nine years. A second spacecraft, Galileo, now in orbit around Jupiter, is equipped with an identical dust detector instrument to Ulysses' dust instrument.  Before and since the Galileo spacecraft's arrival in the Jupiter system in December~1995, investigators recorded more dust stream observations. In July and August 2000, a third spacecraft with a dust detector (combined with a chemical analyzer), Cassini, traveling on its way to Saturn, recorded more high-speed streams of submicron-sized particles from the Jovian system. The many years-long successful Jovian dust streams observations reached a pinnacle on December~30,~2000, when both the Cassini and Galileo dust detectors accomplished a coordinated set of measurements of the Jovian dust streams inside and outside of Jupiter's magnetosphere.  Indirect methods applied by previous researchers have pointed to Io being the simplest explanation for the question of the origin of the Jovian dust streams. We first show by \\underline{direct} methods that Io is the source of the Jovian dust streams. To address the issue of identifying Io directly in the Galileo dust detector data, we apply time-frequency analysis, in particular, Fourier methods, to the Galileo dust data. Additional frequency signatures, such as amplitude modulation, also emerge from the time-frequency analysis.  The second part of this paper focuses on the dust streams dynamics. Here, we apply a detailed Jovian particles and fields model to simulate a dust stream particle's trajectory as the particle moves from Io's orbit through Jupiter's magnetosphere and beyond. Through the model, we show one possible set of parameters that match the travel times seen in the Dec\\-ember~30,~2000 Galileo-Cassini joint dust stream measurements. ",
        "conclusions": ""
    },
    "0206/hep-ph0206046_arXiv.txt": {
        "abstract": " ",
        "introduction": "The mixing of axions with photons and its observable consequences have been analyzed by many authors [1-12]. In the present paper we investigate the changes in the polarization of electromagnetic waves that arise due to its mixing with axions. We are particularly interested in determining if this mixing can explain the polarization anisotropies that have been claimed in Ref. \\cite{Birch,JR,Hutsemekers}.  In Ref. \\cite{Birch,JR} the authors claimed that the observed polarizations from distant radio galaxies and quasars are not isotropically distributed on the dome of the sky. The observable of interest in that study was the angle $\\beta=\\chi-\\psi$, where $\\psi$ is the orientation angle of the axis of the radio galaxy and $\\chi$ the observed polarization angle after the effect of Faraday rotation is taken out of the data. The authors claimed a dipole anisotropy such that the angle $\\beta$ is given by \\begin{equation} \\beta = \\vec\\lambda\\cdot \\hat r \\end{equation} where $\\hat r$ is a unit vector in the direction of the source. The $\\vec\\lambda$ represent the three parameters of this fit. The magnitude of this vector $|\\vec\\lambda|$ is found to be approximately 0.5 and its direction \\begin{equation} \\hat \\lambda = [(0\\ h, 9\\ m)\\pm (1\\ h,0\\ m),-1^o\\pm 15^o]\\ . \\label{axis} \\end{equation} The effect is independent of redshift and was first claimed in Ref. \\cite{Birch} and later verified by a more reliable statistical procedure \\cite{KY} and by compiling a larger data set \\cite{JR}. This anisotropy may be a signal of some local effect arising due to the milky way or the local supercluster. However so far it is not known what physical phenomenon could lead to the observed rotation in polarizations. Within the standard model of elementary particles it is difficult to conceive of a physical mechanism which can lead to this effect. The axion field which arises in many extensions of the standard model of particle physics may provide one possible explanation. However so far it is not known whether such a field can consistently explain this effect. This is one of the questions that we study in the present paper.  Another interesting polarization effect in the electromagnetic waves from distant quasars has been claimed in Ref. \\cite{Hutsemekers,HL}. It was found that optical polarization are aligned  on very large scales. A very striking alignment was found in the region, called A1 in \\cite{HL}, delimited in Right Ascension by $11^{\\rm h}15^{\\rm m} \\le {\\rm RA} \\le 14^{\\rm h}29^{\\rm m}$ and in redshift by $1.0\\le z\\le 2.3$. The polarizations from quasars in any particular spatial region have a tendency to align with one another. The effect was only seen in patches without any evidence of large scale anisotropy. We point out that the center of the A1 region (see Fig. 1 of Ref. \\cite{HL}) is exactly opposite to the axis, Eq. \\ref{axis}, of the anisotropy found in Ref. \\cite{JR}. This might indicate a common origin of these two effects.  Finally we also examine the recent claim \\cite{Csaki} that dimming of distant supernovae \\cite{Perlmutter,Riess} can be explained in terms of axion photon mixing. ",
        "conclusions": "In conclusion we have analyzed the mixing of axions with photons. We have determined how the dispersion relations of the electromagnetic waves are modified due to the presence of axions. We found that for a wide range of parameters this provides the dominant contribution to the changes in polarization of electromagnetic waves due to mixing with axions. We have also determined whether the existence of axions can explain the large scale anisotropies claimed in References \\cite{Birch,JR,Hutsemekers}. We find that the Hutsemekers \\cite{Hutsemekers} effect may be explained by the supercluster magnetic fields if we assume that quasars emit axions such that their flux is of the order of or larger than the photon flux at optical frequencies. The Birch effect \\cite{Birch,JR} may also be explain by assuming axion emission from radio galaxies and quasars but requires a very large flux at radio frequencies. Its explanation in terms of axions is therefore disfavored. Finally we have estimated the contribution of the fluctuations in the plasma density to the dimming of distant supernovae due to axion photon mixing. We find that this effect can be rather large and for a considerable range of currently allowed parameter space can provide an explanation for the supernova dimming.   \\bigskip \\noindent {\\large\\bf Acknowledgements:} We thank John Ralston, Ajit Srivastava and Mahendra Verma for very useful comments. This work is supported in part by a grant from Department of Science and Technology, India."
    },
    "0206/astro-ph0206215_arXiv.txt": {
        "abstract": "We compare deep VLA imaging of the total intensity and linear polarization of the inner jets in the nearby, low-luminosity radio galaxy 3C\\,31 with models of the jets as intrinsically symmetrical, decelerating relativistic flows.  We show that the principal differences in appearance of the main and counter-jets within 30\\,arcsec of the nucleus can result entirely from the effects of relativistic aberration in two symmetrical, antiparallel, axisymmetric, time-stationary relativistic flows.  We develop empirical parameterized models of the jet geometry and the three-dimensional distributions of the velocity, emissivity and magnetic-field structure. We calculate the synchrotron emission by integration through the models, accounting rigorously for relativistic effects and the anisotropy of emission in the rest frame.  The model parameters are optimized by fitting to our 8.4-GHz VLA observations at resolutions of 0.25 and 0.75\\,arcsec FWHM, and the final quality of the fit is extremely good. The novel features of our analysis are that we model the two-dimensional brightness distributions at large number of independent data points rather than using one-dimensional profiles, we allow transverse as well as longitudinal variations of velocity, field and emissivity and we simultaneously fit total intensity and linear polarization.  We conclude that the jets are at $\\approx$52$^\\circ$ to the line of sight, that they decelerate and that they have transverse velocity gradients. Their magnetic field configuration has primarily toroidal and longitudinal components.  The jets may be divided into three distinct parts, based not only on the geometry of their outer isophotes, but also on their kinematics and emissivity distributions: a well-collimated inner region; a flaring region of rapid expansion followed by recollimation and a conical outer region. The inner region is poorly resolved, but is best modelled as the sum of fast (0.8 -- 0.9$c$) and much slower components. The transition between inner and flaring regions marks a discontinuity in the flow where the emissivity increases suddenly.  The on-axis velocity stays fairly constant at $\\approx$0.8$c$ until the end of the flaring region, where it drops abruptly to $\\approx$0.55$c$, thereafter falling more slowly to $\\approx$0.25$c$ at the end of the modelled region.  Throughout the flaring and outer regions, the velocity at the edge of the jet is $\\approx$0.7 of its on-axis value.  The magnetic field in the flaring region is complex, with an essentially isotropic structure at the edge of the jet, but a more ordered toroidal+longitudinal configuration on-axis.  In the outer region, the radial field vanishes and the toroidal component becomes dominant. We show that the emissivity and field structures are inconsistent with simple adiabatic models in the inner and flaring regions. We suggest that the discontinuity between the inner and flaring regions could be associated with a stationary shock structure and that the inferred transverse velocity profiles and field structure in the flaring region support the idea that the jets decelerate by entraining the external medium.  We demonstrate the appearance of our model at other angles to the line of sight and argue that other low-luminosity radio galaxies resemble 3C\\,31 seen at different orientations. ",
        "introduction": "\\label{Introduction}  The flow parameters of jets in extragalactic radio sources have hitherto proven difficult to determine because of the absence of unambiguous diagnostics.  Most progress has been made in the estimation of velocities, particularly where these are thought to be relativistic. The idea that jets in low-luminosity, i.e.\\ FR\\,I \\citep{FR74}, radio galaxies have relativistic speeds rests on five main arguments: \\begin{enumerate}  \\item evidence for relativistic motion on parsec scales in BL Lac objects, coupled with the hypothesis that they are FR\\,I radio galaxies observed at small angles to the line of sight \\citep{UP95};  \\item measurement of apparent proper motions with speeds up to at least $c$ in M\\,87 \\citep*{BZO95};  \\item modelling of relativistic flows, which demonstrates the feasibility of deceleration on kiloparsec scales in realistic galactic atmospheres \\citep*{Bic94, Kom94, BLK96};  \\item the interpretation of correlated depolarization asymmetry and jet sidedness in FR\\,I sources \\citep{Morg97} as a consequence of Doppler beaming and foreground Faraday rotation \\citep{Lai88};  \\item observations of brightness and width asymmetries in FR\\,I jets which decrease with distance from the nucleus and are correlated with fractional core flux, implying that they decelerate and are faster on-axis than at their edges  \\citep{Lai93,Lai96,Hard97,LPdRF}. \\end{enumerate}   This paper reports a detailed study of the applicability of decelerating relativistic jet models to the well-resolved kiloparsec-scale structures in the nearby FR\\,I radio galaxy 3C\\,31.  Our intention is to deconvolve the emission mechanism (synchrotron radiation from a magnetized relativistic flow) from the radio data, without embodying specific preconceptions about the poorly known internal physics. We construct sophisticated three-dimensional models of the effects of relativistic aberration on the appearance of intrinsically symmetric magnetized jets and we fit these models to the observed total and polarized intensity distributions in 3C\\,31.  The aim is to derive robust estimates of the velocity field, the emissivity (combining relativistic particle density and magnetic field strength) and the three-dimensional ordering of the magnetic field (purely geometrical factors independent of its strength). We regard this as a necessary first step: realistic physical models capable of being compared with observations are not yet available.  We are able to reproduce many of the observed features of the jets and we conclude that our models now provide a way to obtain key constraints on the intrinsic properties of extragalactic radio jets.  Section~\\ref{Observations} describes new VLA imaging of the jets in 3C\\,31 that provides a high-quality data set suitable for detailed fitting by our models.  Section~\\ref{Model-theory} first reviews the principles underlying the relativistic jet models, and then goes on to describe how we adjust their parameters  to fit the radio intensity and polarization data.  Section~\\ref{Results} critically discusses the model fits and reviews the main features of the inferred jet velocity field, magnetic structure and emissivity distribution. Section~\\ref{Discussion} discusses more general implications, including: specific problems associated with reproducing the properties of the jets in the region closest to the galactic nucleus; the reasons for the sudden onset of deceleration; evidence for interaction between the jet and the surrounding medium and the applicability of adiabatic models.  Section~\\ref{Angles} illustrates how the jets in 3C\\,31 should appear if orientated at other angles to the line of sight and outlines the applicability of the model to other FR\\,I sources, including those orientated at a small angle to the line of sight.  Section~\\ref{Summary} summarizes our conclusions regarding the kinematics, emissivity and field structure of the three distinct regions of the jet and their implications for future work.  Throughout this paper, we adopt a Hubble constant $H_0$ = 70\\,km s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "\\label{Discussion}  \\subsection{The inner region} \\label{Inner-probs}  Our estimates of the angle to the line of sight and the jet velocity in the inner region are entirely consistent with those of \\citet{Lara97} for the jets on parsec scales.  For our preferred value of $\\theta$, their velocity range is $0.81 \\leq \\beta \\leq 0.998$, consistent with our central velocity for the inner region ($\\beta \\approx 0.87$), but also allowing significant deceleration from parsec to kiloparsec scales.  The inner region poses a problem for any decelerating-jet model, in which the jet-to-counterjet intensity ratio (sidedness) must decrease with distance from the nucleus.  We would expect the sidedness ratio to have a maximum in the inner region, in which case the counter-jet would be invisible.  In fact, the inner region clearly has a {\\em lower} sidedness, on average, than that in the flaring region.  The brightness distribution is dominated by a few knots, so one possibility is that we are being fooled by stochastic variations.  Alternatively, there could be a small amount of relatively slow-moving material in the shear layer, surrounding a fast spine.  If a very slow component also exists further out, it cannot have a noticeable effect on the brightness distribution: the fact that the sidedness ratio at the edges of the jet in the flaring region differs significantly from unity requires that the emissivity of any very slow component becomes insignificant on large scales.  This component therefore has a negligible effect on the fits beyond the inner jet, and its properties are constrained only by the intensity fits in the inner region.  The faster component of the flow that dominates the outer jets is relatively faint in the inner region (because its Doppler factor and intrinsic emissivity are both lower than on larger scales), so the modelling of the inner region is almost decoupled from that of the rest of the jets in the spine/shear-layer fits.  In the best fit SSL model, the slow component has been introduced by assigning a velocity of 0.06$c$ to the shear layer while the spine has a velocity of 0.87$c$, i.e.\\ we have allowed the slow component in the inner region to substitute for the faster-moving shear layer component that is required to explain the flaring region via an unphysical jump condition at the boundary. We emphasize that higher resolution imaging of the inner jets is needed to obtain firmer constraints on the transverse velocity distribution in this region, and to explore how this distribution evolves as the jet enters the flaring region.  Our best guess at the geometry is sketched in Fig.~\\ref{inner-sketch}.  If this picture is correct, there must still be an increase in emissivity at the flaring point, but the values of $g$ quoted in Table~\\ref{Params} will be inaccurate. Finally, we note that an additional component with $\\beta$ significantly higher than 0.9 would be severely Doppler-dimmed even in the approaching jet, and therefore very difficult to detect.  \\begin{figure} \\epsfxsize=8cm \\epsffile{fig21.eps} \\caption{ Grey-scale images of the rms magnitudes of the magnetic field components as fractions of the total field for the SSL model.  Top: radial component $\\langle B_r^2 \\rangle^{1/2} / B$; centre: toroidal component $\\langle B_t^2 \\rangle^{1/2} / B$; bottom: longitudinal component $\\langle B_l^2 \\rangle^{1/2} / B$.  The arrows marks the location at the edge of the jet where the three field components are roughly equal, as discussed in the text.  The radial component is constrained to be zero for the spine and the $s = 0$ streamline in the shear layer.  The values for the inner region are poorly determined (Table~\\ref{Params}). \\label{Bcomp-grey} } \\end{figure}  \\subsection{The onset of flaring and deceleration} \\label{Flaring}  We have shown that the onset of deceleration is marked by a large increase in rest-frame emissivity and a major change in the jet collimation.  It is not merely that the jet becomes gradually brighter as it decelerates and Doppler suppression is reduced: there is also a discontinuity at the inner boundary.  One possibility is that the jet is supersonic, over-pressured and expanding freely in the inner region.  In that case, the internal pressure would fall until it drops below that of the external medium, at which point a reconfinement shock forms \\citep{Sand83}.  The reconfinement shock is followed by a second shock at which the jet becomes {\\em overpressured} with respect to the external medium and it this feature which is most plausibly identified with the flaring point.  For a relativistic jet, \\citet{Kom94} shows that the shock forms at a distance \\[ z_{\\rm shock} \\approx \\left ( \\frac{2 \\Phi}{3 \\pi p_{\\rm ext} c} \\right )^{1/2}\\] where $\\Phi$ is the energy flux through the jet and $p_{\\rm ext} \\approx 3 \\times 10^{-11}$\\,Pa \\citep{Hard} is the external pressure. This would be consistent with the observed inner boundary distance of 1.1\\,kpc for an energy flux of $\\approx 5 \\times 10^{37}$\\,W, somewhat higher than that the value of $\\approx 1 \\times 10^{37}$\\,W estimated by \\citet{LB02} from a conservation-law analysis.  We see no evidence for any simple shock structure at the inner boundary, although the emission there is not completely resolved and there are obvious (non-axisymmetric) knots at the beginning of the flaring region. If the inner region is in free expansion, we can estimate the initial Mach number of the flow from the opening angle: $\\arctan(\\xi_{\\rm i}) \\approx {\\cal M}$ where ${\\cal M} = (\\Gamma\\beta)/(\\Gamma_s\\beta_s)$ is the generalized Mach number defined by \\citet{Kon80}, $\\beta_s = c_s/c$, $c_s$ is the internal sound speed, and $\\Gamma_s = (1-\\beta_s^2)^{-1/2}$.  The observed value of $\\xi_{\\rm i} = 6.7^\\circ$ corresponds to ${\\cal M} \\approx 8.5$ and hence to $\\Gamma \\approx 6.1$ if the inner jet has the sound speed $c_s = c/\\sqrt 3$ of an ultra-relativistic gas.  This initial velocity is considerably faster than we have inferred for the inner region but, as mentioned in Section~\\ref{Inner-probs}, we cannot exclude the presence of such higher-velocity material there.  A second possibility which has frequently been discussed in the literature is that the flaring point marks the onset of turbulence, or the position at which Kelvin-Helmholtz instabilities become non-linear (e.g.\\ \\citealt{Baa80,Beg82,Bic84,Bic86,DeY96,RHCJ99,RH00}).  We will show elsewhere \\citep{LB02} that conservation-law analysis favours the hypothesis that the flaring point is associated with a stationary shock, primarily because it suggests that the jet is over-pressured at the beginning of the flaring region. This does not, of course, exclude the subsequent development of entrainment (and presumably turbulence), as we now discuss.   \\begin{figure} \\epsfxsize=7cm \\epsffile{fig22.eps} \\caption{A sketch of one possible geometry for the inner region and its transition to the flaring region, incorporating a slow boundary layer which does not persist at large distances from the nucleus.\\label{inner-sketch}} \\end{figure}  \\subsection{Evidence for interaction with the surrounding medium} \\label{Interaction}  It is generally accepted that jets in FR\\,I radio galaxies decelerate by picking up matter, but it is by no means clear whether the principal source of additional material is mass loss from stars \\citep{Phi83,Kom94,BLK96} or entrainment across the jet boundary \\citep{Baa80,Bic84,Bic86,DeY96}: both are expected to be important.  Our models require significant transverse velocity gradients, in the sense that the edge of the jet is travelling about 30\\% more slowly than the centre.  These gradients are prima facie evidence for interaction between the flow and the external medium.  There is no reason why mass input from stars should generate such gradients \\citep{BLK96}, although a pre-existing gradient might be preserved as a jet becomes mass-loaded. The {\\em form} of the transverse velocity profile in our best-fitting models varies surprisingly little as the jet decelerates, but the error analysis of Section~\\ref{Tolerances} shows that the situation might be more complicated: a top-hat velocity profile at the inner boundary is consistent with the data, so the profile could still evolve significantly along the jet.  The presence of large quantities of very slow material at the edges of the flaring and outer regions is firmly excluded, however (Table~\\ref{Params}).  A second piece of evidence favouring deceleration by interaction with the external medium is the complex field structure in the flaring region, where we were forced to introduce a significant radial component, increasing towards the edge of the jet, in order to explain the low degree of linear polarization.  This radial field component would not be expected from simple passive evolution of a mixture of longitudinal and toroidal field in the smooth velocity field we assume.  The most natural way to generate such a radial field component is for the flow to have a disordered, turbulent character towards the jet edges such as might result from large-scale eddies.  This is precisely the situation expected at the edge of the jet in the initial ``ingestion'' phase of the entrainment process \\citep{DeY96}.  The velocity field is then likely to have significant small-scale structure which is not included in our model, but our estimates of average bulk flow speed are unlikely to be seriously affected.  Even if there is no dissipation or dynamo action in such a turbulent flow, there will be significant amplification of the magnetic field by shear, so the simplest adiabatic models, which assume laminar flow (Section~\\ref{Phys-parms}), will be inappropriate.  Another way to distinguish stellar mass loading from entrainment across the jet boundary is to ask whether stellar processes can provide the mass input rate required to produce the observed deceleration.  It is clear from the work of \\citet{Kom94} and \\citet{BLK96} that a jet which is decelerated purely by stellar mass loading will tend to reaccelerate on large scales, where the stellar density becomes low but the outward pressure gradient and buoyancy force are still appreciable. Our models require continuous deceleration in the outer region, favouring boundary-layer entrainment as the dominant mechanism there.   We address this question via a conservation-law analysis in \\citet{LB02}, where we conclude that entrainment dominates after the beginning of the flaring region.  Little is known about the properties of turbulent relativistic shear layers, or of the viscosity mechanisms likely to predominate in magnetized relativistic jets.  We cannot therefore relate the deduced velocity profiles to the internal physics of the jets.  We note however that \\citet{Baa80} computed steady-state models for viscous jets in constant-pressure atmospheres and estimated both the transverse velocity profiles and appearance of the jets (on the assumption that the emissivity is directly proportional to the viscous dissipation) for several forms of the viscosity.  Baan's models generally predicted extended low-velocity wings that do not match our derived profiles.  He did however discuss circumstances under which flat-topped velocity profiles such as those inferred here might arise, including that of an electron-positron jet.  \\subsection{The emissivity profile and adiabatic models} \\label{Phys-parms}  We have determined the variation of $n_0 B^{1+\\alpha}$ (proportional to the emissivity) in the rest frame of the emitting material.  Separation of this variation into particle and field contributions requires additional assumptions.  The X-ray emission from the jets \\citep{Hard} is most likely to be synchrotron, rather than inverse Compton radiation, so we cannot use it to decouple the particle and field components.  We therefore postpone a discussion of the variation of pressure and density along the jet to \\citet{LB02}, where we also consider X-ray observations of the surrounding hot gas.  A number of authors \\citep{Bau97,Fer99,Bondi00} have recently re-opened the possibility that the jets in FR\\,I radio galaxies are adiabatic in the sense first defined by \\citet{Bur79}, i.e.\\ that the particles suffer only adiabatic energy losses, there are no dissipative processes causing particle acceleration or field amplification and the magnetic field is convected passively with the flow. We defer a full discussion of this question to a later paper, since our data and models are both substantially more complicated than is allowed by the analytical approaches in the literature \\citep{Bau97}.  The simplest adiabatic models do not allow for any turbulent flow (Section~\\ref{Interaction}) and there is independent evidence for particle acceleration in 3C\\,31's jets from X-ray observations \\citep{Hard}.  Nevertheless, we can make a number of preliminary qualitative points.  We take the analytical formulae from \\citet{Bau97}.  In the absence of velocity shear and in the quasi-one-dimensional approximation, the field components vary as: \\begin{eqnarray*} B_l & \\propto & x^{-2} \\\\ B_t & \\propto & (x\\beta\\Gamma)^{-1} \\\\ B_r & \\propto & (x\\beta\\Gamma)^{-1} \\\\ \\end{eqnarray*} where $x$ is the jet radius. For a purely longitudinal field, this leads to a variation of the rest-frame emissivity: \\begin{eqnarray*} \\epsilon & \\propto & (\\Gamma\\beta)^{-(2\\alpha+3)/3}x^{-(10\\alpha+12)/3} \\\\ & = & (\\Gamma\\beta)^{-1.37}x^{-5.83} \\\\ \\end{eqnarray*} and for a perpendicular field ($B_l = 0$): \\begin{eqnarray*} \\epsilon & \\propto &  (\\Gamma\\beta)^{-(5\\alpha+6)/3}x^{-(7 \\alpha+ 9)/3} \\\\ & = & (\\Gamma\\beta)^{-2.92}x^{-4.28} \\\\ \\end{eqnarray*}  The inner region poses a severe problem for the simplest adiabatic models: we have no evidence for deceleration so, if the conical region is fully filled, we would expect a very rapid brightness decline away from the nucleus ($\\propto z^{-4.28}$ on-axis even in the perpendicular-field case) compared with our estimates of $\\propto z^{-1.96}$ for the spine and $ \\propto z^{-1.33}$ for the shear layer. Even for the steepest emissivity fall-off allowed by our error analysis (Table~\\ref{Params}), the indices are grossly discrepant. We have already argued that much of the emission in the inner region may come from a surface layer (Section~\\ref{Inner-probs}) and the assumption that the radiating material expands with constant opening angle may be invalid.  For the flaring and outer regions, we have computed the emissivity variations for the parallel- and perpendicular-field cases using our model for the radius and velocity of the jet. The results are shown in Fig.~\\ref{Em-long}, where we have normalized the adiabatic models to the observed emissivities at the beginnings of the flaring and outer regions. Two example streamlines are shown for the SSL model: on-axis in the spine and at the inner edge of the shear layer.  The adiabatic models predict emissivities which fall far more rapidly than is observed in the flaring region: the deceleration is too little and too late to compensate for the expansion. In the outer region, by contrast, the perpendicular-field adiabatic model predicts emissivities fairly close to those observed.  \\begin{figure*} \\epsfxsize=12cm \\epsffile{fig23.eps} \\caption{ The best-fitting model of the 3C\\,31 jets viewed at various angles $\\theta$ to the line of sight with a beam of 0.75\\,arcsec FWHM. Left panel: logarithmic contours with fixed sensitivity, i.e.\\ with the same lowest contour in all plots.  Right panel: logarithmic contours with fixed 750:1 dynamic range i.e.\\ with the same percentages of the peak intensity in all plots.  Both sets of plots cover $\\pm$27\\,arcsec from the nucleus and the angular scale is indicated by the bar at the top of the diagram. \\label{otherangles} } \\end{figure*}  There are two other fundamental problems with the adiabatic models. First, the field structure in the flaring region is not consistent with passive convection in a smooth, axisymmetric velocity field.  Our assumed velocity field acts so as to shear an existing radial component, thereby amplifying the component along the flow.  It cannot, therefore, create the region of approximately isotropic field at the edge of the flaring region starting with what is essentially a mixture of toroidal and longitudinal components.  Second, the assumed velocity field {\\em cannot} change the ratio of radial to toroidal field.  It is clear from Fig.~\\ref{Bcomp-grey} that the radial component essentially disappears at some point after the flaring region.  We conclude that simple adiabatic models could not describe the inner and flaring regions, even if more realistic field configurations and the effects of velocity shear were to be included, but that a model of this type may apply to the outer region, at least if the radial field component is mostly eliminated by the outer boundary.  Further investigation of this set of problems is outside the scope of the present paper and will be presented elsewhere.   \\label{Summary}  \\subsection{Method}  We have shown that an intrinsically symmetrical, decelerating relativistic jet model containing simple prescriptions for the velocity field and emissivity with a locally random but anisotropic magnetic field, accounts for the major features revealed by deep VLA imaging of the straight segments of the jet and counter-jet in 3C\\,31.  The principal new features of our approach are: \\begin{enumerate} \\item the use of three-dimensional (but axisymmetric) parameterized models of velocity, emissivity and field ordering; \\item rigorous calculation of synchrotron emission, including both relativistic aberration and anisotropy in the rest frame; \\item fitting to images with many independent data-points in linear polarization as well as total intensity using a robust optimization algorithm. \\end{enumerate}  \\subsection{Principal regions of the jets} \\label{Regions}  A major result of this modelling is that the three regions of the jet that were initially identified purely from the shape of the outer isophotes (Fig.~\\ref{Geom-sketch}) are also regions with distinctly different internal variations of velocity and emissivity\\footnote{All distances in this section are measured in a plane containing the jet axis, corrected for projection using the angle to the line of sight for the best-fit SSL model.}.  \\subsubsection{The inner region (0 to 1.1\\,kpc)}  Our conclusions for this region are tentative because of the limited transverse resolution of the data.  The region is characterized by: \\begin{enumerate} \\item low intrinsic emissivity; \\item slow lateral expansion (a cone of intrinsic half-angle 6.7$^\\circ$) and \\item a significant component of emission arising in slow-moving material. \\end{enumerate} The fitted central velocity is 0.8 -- 0.9$c$.  We have no evidence for deceleration in this region, but we cannot exclude the presence of higher-velocity (Doppler-hidden) flow components.   Simple adiabatic models are grossly inconsistent with the emissivity profile.  \\subsubsection{Flaring region (1.1 to 3.5\\,kpc)}  This region was defined initially by the more rapid spreading of its outer isophotes.  Our modelling shows it to be a region in which several dramatic changes in the other jet characteristics occur together: \\begin{enumerate} \\item The jets decelerate rapidly to an on-axis velocity of 0.55$c$ after an initial slow decline from 0.77$c$. \\item They maintain a transverse velocity profile in which the edge velocity drops to approximately 70\\% of the on-axis value. \\item The intrinsic emissivity increases abruptly at the boundary with the inner region, then declines with distance from the nucleus, $z$, as $z^{-3.1}$ in the shear layer and $z^{-2.5}$ in the spine. \\item The emissivity at the edges of the jet drops to about 20\\% of that on the jet axis. \\item The radial component of the magnetic field in the shear layer becomes significant, rising from zero at the spine boundary to 90\\% of the toroidal and longitudinal components at the outer edge of the layer, i.e.\\ the field is essentially isotropic at the outer boundary of the shear layer in this region. \\item The ratio of longitudinal to toroidal field strength decreases slightly from about 1.1 to 0.8, independent of radius in the jet. \\end{enumerate} The sudden increase in rest-frame emissivity at the flaring point suggests that there is a discontinuity in the flow, perhaps a stationary reconfinement shock system.  The brightness and polarization structure in this region cannot be described by a simple adiabatic model.  The transverse velocity profile and the growth of the radial field component strongly suggest that entrainment across the jet boundary becomes important.  \\subsubsection{Outer region (3.5 to 12\\,kpc)}  In this region, the jets continue to expand on a cone of intrinsic half-angle 13.1$^\\circ$. \\begin{enumerate} \\item The jets decelerate less rapidly, reaching an on-axis velocity of 0.26$c$ by 10\\,kpc. \\item The intrinsic emissivity in the shear layer declines more slowly ($\\propto z^{-1.4}$) with distance from the nucleus. \\item The transverse velocity and emissivity profiles remain essentially unchanged from those in the flaring region. \\item The ratio of radial to toroidal magnetic field strength decreases, becoming $<20$\\% throughout the jet by 10\\,kpc. \\item The ratio of longitudinal to toroidal magnetic field in the shear layer continues to decrease, from 0.8 to 0.5 by 10\\,kpc. \\end{enumerate} Although the emissivity fall-off is much closer to that predicted by a perpendicular-field, laminar adiabatic model, more work is needed to test this idea for realistic field and particle distributions.  Beyond the end of the outer region, intrinsic environmental asymmetries begin to dominate, as evidenced by the large-scale bending of both jets.  \\subsection{Implications for unified models}  We have also calculated the change in appearance of our model brightness and polarization distributions as functions of orientation.  These are in good qualitative agreement with observations of other well-observed jets and we therefore expect the model (with some parameter variations) to apply to FR\\,I jets in general.  Figures~\\ref{otherangles} and \\ref{otherangles-hires} show that the intensity changes are considerably more complex than would be expected for single-velocity jets. They emphasize the need for high dynamic range and sensitivity to possible wide-angle jet structures when assessing whether observed jet properties are consistent with unified models.    We predict changes in polarization with orientation (Figure~\\ref{otherangles-pol}): these provide an independent test of unified models provided that our proposed field configuration is present in all FR\\,I jets.  \\subsection{Further work}  We now intend to model other resolvable bright jets in FR\\,I radio galaxies to determine the extent to which their observed brightness and polarization properties resemble those of 3C\\,31. We expect to be able to infer their velocity, emissivity and magnetic-field distributions, building on the broad success of the jet-deceleration model in accounting for the statistical asymmetries of the B2 sample of FR\\,I sources \\citep{LPdRF}.  Other sources showing well-collimated inner jets and rapid flaring include NGC\\,315 \\citep{Vent}, PKS\\,1333$-$33 \\citep*{KBE} and 3C\\,449 \\citep{Fer99}, and it seems likely that the regimes of collimation behaviour we have identified in 3C\\,31 are common in FR\\,I sources.  We aim to study a sample of sources with a range of angles to the line of sight, if possible distributed isotropically, in order to test the results of Section~\\ref{Angles}.  We also plan to develop a more sophisticated error analysis in order to assess confidence levels with some degree of rigour.  3C\\,31 has been cited as the archetypal FR\\,I source, but is actually in the minority in having diffuse ``tails'' of emission extending to large distances from the core rather than confined bridges analogous to the lobes of FR\\,II sources \\citep{DeR}. Significant differences in dynamics (especially entrainment of the surrounding medium) might be expected between the two classes.  We also expect that the deceleration process should depend on the jet power and the external environment.  In \\citet{LB02}, we present a dynamical model for the jets in 3C\\,31, based on the velocity field derived in the present paper, a description of the surrounding galactic atmosphere derived from {\\em Chandra} and ROSAT observations \\citep{Hard} and application of conservation laws following \\citet{Bic94}. This approach should also be extensible to other sources. Our results favour entrainment across the boundary layer as the origin of the majority of the mass-loading of the jets in 3C\\,31, but it will be important to explore this in other large-scale FR\\,I radio galaxies.  We should seek further evidence for the entrainment process, such as the reduced polarization near the boundaries of the flaring regions.  Our ultimate goal is to replace the empirical descriptions of velocity, emissivity and field structure with realistic physical models. Although this is some way off, we have developed a self-consistent adiabatic model which can handle arbitrary field configurations and (laminar) velocity fields in a relativistic jet, with the aim of establishing whether any of the flow regions we have identified can be described in this way.  If our interpretation of the emission from the inner region of 3C\\,31's jets is correct, observations of the apparent brightness and motions of FR\\,I jets on even smaller scales will {\\it not} be sensitive to the properties of the underlying bulk flow, but only to those of its slowest-moving components, which may be essentially stochastic.  Improved transverse resolution of the inner jets in such sources will be required to determine the origin and distribution of the slow-moving material, and the extent to which these innermost regions of FR\\,I jets resemble the larger-scale jets in FR\\,II sources, e.g.\\ those in 3C\\,353 \\citep*{Swa98}.  This will require greater sensitivity and longer baselines than are currently available with the VLA or MERLIN.  Finally, a number of FR\\,I sources (including 3C\\,31) have been detected at X-ray and/or optical wavelengths (e.g. \\citealt{Hard66B,Hard,Wor,Sparks,CenA,M87opt,M87X}). The radiation is most plausibly produced by the synchrotron process over the entire observed frequency range, and the shape of the spectrum therefore carries information about particle acceleration and energy loss. It will be important to incorporate descriptions of these processes into our models."
    },
    "0206/astro-ph0206509_arXiv.txt": {
        "abstract": "Radio plasma injected by active radio galaxies into clusters of galaxies quickly becomes invisible due to radiative losses of the relativistic electrons. In this talk, the fate of radio plasma and its role for the galaxy cluster is discussed: buoyancy removes it from the central regions and allows to transfer its energy into the ambient gas. The remaining low energy electron populations are still able to emit a low luminosity glow of observable radiation via synchrotron-self Comptonized emission. Shock waves in the ambient gas can re-ignite the radio emission. ",
        "introduction": "The jets of powerful radio galaxies inflate large cavities in the intracluster medium (ICM) that are filled with relativistic particles and magnetic fields. Synchrotron emission at radio frequencies reveals the presence of electrons with several 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' ({En{\\ss}lin} 1999). Their existence as a separate component of the ICM is supported by the detections of cavities in the X-ray emitting galaxy cluster gas ({B{\\\"o}hringer} et al. 1993; {Carilli}, {Perley}, \\& {Harris} 1994; {Huang}, \\& {Sarazin} 1998; {McNamara} et al. 2000; {Fabian} et al. 2000; {Finoguenov} \\& {Jones} 2001; {Fabian} 2001; {McNamara} 2000; {Heinz} et al. 2001; {Schindler} et al. 2001; and others). 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 (Churazov et al. 2001). 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.  \\begin{figure}[t] \\plotone{fig1.eps} \\caption{Bubble's central X-ray contrast (compared to the undisturbed cluster) for various angles between plane of sky and Bubble's trajectory, its radio flux, and its rising time as a function of the (unprojected) radial position. An X-ray background with 1/100 of the central cluster surface brightness is assumed, which is responsible for the strong decrease at large radii of the X-ray contrast in the $0^\\circ$ and $45^\\circ$ cases.} \\label{eps5} \\end{figure} ",
        "conclusions": ""
    },
    "0206/astro-ph0206353_arXiv.txt": {
        "abstract": "We report on the discovery of a free-floating methane dwarf toward the direction of the young star cluster $\\sigma$\\,Orionis. Based on the object's far-red optical and near-infrared photometry and spectroscopy, we conclude that it is a possible member of this association. We have named it as S\\,Ori\\,J053810.1--023626 (S\\,Ori\\,70 is the abridged name). If it is a true member of $\\sigma$\\,Orionis, the comparison of the photometric and spectroscopic properties of S\\,Ori\\,70 with state-of-the-art evolutionary models yields a mass of 3\\,$^{+5}_{-1}$ Jupiter mass for ages between 1\\,Myr and 8\\,Myr. The presence of such a low-mass object in our small search area (55.4\\,arcmin$^2$) would indicate a rising substellar initial mass function in the $\\sigma$ Orionis cluster even for planetary masses. ",
        "introduction": "Since the discovery of brown dwarfs (i.e., objects unable to burn hydrogen stably in their interiors and with masses below 72\\Mj; Kumar \\cite{kumar63}; Chabrier et al$.$ \\cite{chabrier00}) both in the field and in young open clusters (see Basri \\cite{basri00} for a review), many questions remain unsolved. A very important one is the minimum mass for the formation of very low mass objects in isolation, which would represent the bottom end of the initial mass function (IMF) for free-floating objects. B\\'ejar et al$.$ \\cite{bejar01} have scrutinized the substellar population of the young, nearby $\\sigma$ Orionis cluster, and derived a rising IMF (d$N$/d$M$\\,$\\sim$\\,$M^{-0.8\\pm0.4}$), which is complete for brown dwarf masses in the range of 72--13\\Mj. Other authors find similar substellar IMFs (e.g., Luhman et al$.$ \\cite{luhman00}). Very recent photometric and spectroscopic searches (Lucas \\& Roche \\cite{lucas00}; Zapatero Osorio et al$.$ \\cite{osorio00}) suggest that the IMF extends further below the deuterium burning mass threshold at around 13\\Mj~(Saumon et al$.$ \\cite{saumon96}; Burrows et al$.$ \\cite{burrows97}). We will refer to this mass regime as the ``planetary mass'' domain. The least-massive objects so far identified in young stellar clusters of Orion have masses around 5--10\\Mj~(Lucas et al$.$ \\cite{lucas01}; Mart\\'\\i n et al$.$ \\cite{martin01}), and cover the full range of the spectral type L. For the formation scenarios of such objects it would be very important to see down to which masses they can be found in isolation.  Here we report on the discovery of a methane dwarf toward the direction of Orion. In this paper we present evidence for its membership in the $\\sigma$\\,Orionis star cluster, which implies that this object is likely the least massive planetary mass body imaged to date outside the solar system. Cluster membership is discussed in Section \\ref{membership}, and mass determination based on state-of-the-art evolutionary models and further discussion are presented in Section \\ref{mass}. ",
        "conclusions": ""
    },
    "0206/astro-ph0206165_arXiv.txt": {
        "abstract": "{We investigate the star cluster system in the starburst galaxy NGC~7673 using archival {\\it Hubble Space Telescope} WFPC2 broad-band images. For the first time we are able to examine the internal structures of the prominent optical clumps in this galaxy. The clumps are composed of young stars, $16-33 \\%$ of which are in bright star clusters. We identify 268 star cluster candidates in both the F555W and F814W images, and 50 clusters with the F255W filter. These data allow us to estimate ages and masses using color-magnitude and two-color diagrams for our sample. We find a population of young, $< 6$~Myr clusters located throughout the galaxy with concentrations in the clumps. Star cluster mass estimates are $5-50 \\times 10^{4}$~M$_{\\odot}$ for the brightest objects. The starburst remains active in physically well-separated regions, indicating a widespread starburst trigger. We discuss clump lifetimes, their implications for the future evolution of NGC~7673, and possible relationships to high redshift starbursts.} ",
        "introduction": "Starburst galaxies churn gas into stellar light at tremendous rates, so fast that the time to exhaust this material is short compared to the age of the universe (e.g., \\cite{Mira92}).  In recent years they have also become notorious for their propensity to form swarms of star clusters, including massive, compact super star clusters (SSCs), which have been suggested by many authors to be the progenitors of globular clusters (e.g., \\cite{HOG94}, \\cite{WS95}, \\cite{AZ01}).  The starburst phenomenon occurs rarely in the local universe, but increases in frequency at larger lookback times (e.g., \\cite{H77}, \\cite{Liu98}, \\cite{Glaze99} and references therein), which may be connected to the observed excess of distant ``faint blue galaxies'' (\\cite{E97}).  Studies of moderate redshift ($z \\sim 1$) galaxy samples reveal substantial populations of luminous galaxies with high star formation rates (SFRs) (\\cite{CHS95}, \\cite{Oetal01}).  These include the blue ``compact narrow emission line galaxies'' (CNELGs; \\cite{Ketal94}, \\cite{Ketal95}, Guzm\\'{a}n et al. 1996, 1997, 1998, \\cite{Petal97}).  Similar objects have been categorized into various groups with various names, such as the ``luminous blue compact galaxies'', or LBCGs (our preferred term for such objects; \\cite{Jetal01}). Galaxies of this class have small linear sizes ($<$20~kpc), luminosities near or above $L^*$, high surface brightnesses, strong emission lines, and blue colors, all features that are indicative of enhanced SFRs.  The evolutionary fate of moderate redshift LBCGs is a subject of some debate, and depends strongly on the current and future SFRs in comparison with the extent of their gas reservoirs.  One scenario holds that the most compact of the LBCGs will evolve into spheroidal stellar systems covering a range in mass (e.g., \\cite{BR92}, \\cite{Getal98}).  The similarity of this subclass of LBCGs to luminous, young, star-forming \\ion{H}{ii} galaxies has led to a classification of some systems as ``\\ion{H}{ii}-like'' (\\cite{Petal97}, \\cite{Getal97}).  Passive evolutionary models predict that after 4-6 Gyr, the luminosities and surface brightnesses of these \\ion{H}{ii}-like LBCGs could resemble local low-mass ellipticals, like NGC~205 (\\cite{Getal98}). The viability of this scenario depends on the star formation in LBCGs having a duration of less than 1~Gyr, and in that time they must lose almost all their gas (\\cite{Petal01}), which is difficult to accomplish in any but the least massive dwarfs (\\cite{MLF99}, \\cite{FT00}).  Alternatively, LBCGs could be analogs to intense starbursts in nearby disk galaxies, especially when observed in near face-on orientations that favor the escape of blue/ultraviolet light from the starburst region (\\cite{GHB89}, \\cite{GCH00}, \\cite{BvZ01}).  The evolutionary scenario for disk-like LBCGs does not necessarily lead to their descendants being spheroidal galaxies.  In nearby examples of this starburst mode, relatively minor interactions between galaxies can yield major starbursts in which the disk is only moderately perturbed (e.g., M82).  If the post-burst systems have sufficient gas supplies to allow star formation to continue, albeit at lower rates than in the current bursts, then they can appear at the present epoch as comparatively normal star-forming disk galaxies (\\cite{Petal97}, \\cite{Hetal00}).  Related to these evolutionary questions is the issue of bulge formation.  Theoretical work has shown that gas-rich ($> 10 \\%$ in \\ion{H}{i}) disks subject to gravitational instability may form large clumps of gas, up to $10^{9}$ M$_{\\odot}$, which rotate in the plane of the disk (\\cite{EKT93}, \\cite{N99}). Such clumps experience strong star formation, resulting in a morphologically peculiar galaxy, and could appear like the 'clumpy' galaxy investigated here.  The clumps suffer dynamical friction from the surrounding visible and dark matter, leading them to spiral inwards and accumulate in the central region, potentially forming a bulge.  In this scenario the disk forms {\\it first}, hosts the initial round of star formation, produces a bulge. The remaining gas could then fuel a more sedate course of evolution. However, a key factor in this model is whether the clumps can retain their identities over the $\\sim$100~Myr time scales required for dynamical friction to act (\\cite{N99}).  To clarify the general issues of how starbursts connect to star formation processes and galaxy evolution, we concentrate on understanding nearby starbursts whose internal properties are observable.  With this goal in mind, the Wide Field Planetary Camera 2 (WFPC2) Investigation Definition Team GTO program included exploration of small scale structures in blue starburst galaxies with M$_{B} < -18$ and high optical surface brightness, many of which are members of the LBCG class (\\cite{GHC00}).  NGC~7673 (a.k.a. IV Zw 149, Markarian 325, UGC 12607, and UCM 2325+2318), a 'disturbed spiral' or 'clumpy irregular' galaxy was chosen for a multi-wavelength imaging investigation extending from the mid-ultraviolet to the near infrared. NGC~7673 is a member of the LBCG class, and a luminous infrared source (e.g., \\cite{GHB89}, \\cite{Hetal89}, \\cite{SM96}).  As such, it presents a unique opportunity to understand the inner workings of a high SFR object and its potential evolutionary connections to distant and present-day galaxy populations.  This paper presents an analysis of star clusters in the clumpy starburst regions of NGC~7673 based on archival images obtained with WFPC2 on the {\\it Hubble Space Telescope}. We assume H$_{0} = 70$~km~s$^{-1}$~Mpc$^{-1}$; the recession velocity of 3408~km~s$^{-1}$ from the NASA/IPAC Extragalactic Database\\footnote {The NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.} for NGC~7673 then implies a distance of 49~Mpc, and a projected scale of 250 pc per arcsec. The distance modulus for this galaxy is 33.45 and one WFC pixel covers a projected scale of approximately 25 pc.  The next section summarizes properties of this unusual galaxy.  Our observations and photometry of star clusters are detailed in \\S 3, and \\S 4 provides information on our choice of models for the spectral evolution of the clusters.  In \\S 5 we discuss the morphology of NGC~7673 as seen with WFPC2, \\S 6 discusses the star cluster colors and magnitudes in terms of their ages, and \\S 7 briefly characterizes the star-forming clumps. Our results are discussed \\S 8, and in \\S 9 we present our conclusions. ",
        "conclusions": "NGC~7673 is a fascinatingly disturbed galaxy, with a total SFR ($\\approx 10-20$~M$_{\\odot}$~yr$^{-1}$) and a SF intensity of for Clump~B ($\\sim 1$~M$_{\\odot}$yr$^{-1}$kpc$^{-2}$) that put it in the range populated by Lyman break galaxies (\\cite{Petetal01}). Due to its combination of proximity for a galaxy with CNELG properties and favorable face-on orientation, it affords an excellent opportunity to explore the structure of a major starburst. While further study is needed to fully disentangle the effects of age and reddening on star cluster colors, our initial investigation reveals several interesting trends.  1. Working from WFPC2 images, we have identified 50 star cluster candidates where we measured F255W, F555W and F814W magnitudes, and 268 with F555W and F814W photometry. While we find a broad range of star cluster colors, the bluer and brighter clusters are strongly concentrated into the main starburst `clumps'.  2. In our three color sample we mainly measure clusters with ages of $<$20~Myr. All of the main optical clumps, A, B, C, D, and F contain star clusters with ages of $5$~Myr or younger. However, the nuclear region, clump A, may contain older clusters than are found in clumps B and C. Even so we do not find an obvious age ordering of the clumps. Further observations with more filters and deeper data are needed to allow better separation of clusters which are red due to age versus young clusters reddened by interstellar dust.  3. The four brightest clumps are composed of $> 50$ star clusters detected from the UV through the optical, providing $\\sim16-33\\%$ of the luminosity from the clump as a whole. These results are consistent with other estimates for the fraction of light contributed by compact star clusters in starbursts (e.g., \\cite{Metal95})  Clump~F is distinguished by not being an association of star clusters. It is instead dominated by a single object, for which we derive an age of $4-5$~Myr and a mass of $3 \\times 10^{5}$~M$_{\\odot}$. If the age is 4~Myr or slightly less, it should show the presence of Wolf-Rayet stars in its spectrum. Perhaps Clump~F resembles a more luminous version of compact young stellar complexes with a dominant super star cluster, such as the Hodge complex in NGC~6946 (\\cite{Laretal02}) or the extended knot S that is composed of luminous young stars in the Antennae (\\cite{Wetal99}).  4. The star formation process in NGC~7673 differs from that seen in normal spirals in that the disk has broken up into relatively well-defined clumps that contain most of the current star-forming activity. The clumps may be understood as the result of large spatial scale instabilities in perturbed gaseous disks, as suggested by Elmegreen et al. (1993). However, it is not yet clear if the current clumps can survive long enough for dynamical friction to act to bring clumps into the center of the galaxy where they can form a bulge, as suggested by Noguchi (1999).  The unusual structure of NGC~7673 and other nearby starburst galaxies indicate that the impact of highly coeval populations of star clusters should be taken into account when considering how best to model galaxies with intense star formation. These issues extend from implications for the integrated rest-frame UV spectra to their impact on the ISM, and even to the subsequent evolution of the galactic disks (\\cite{K02b}). Furthermore, we must take into account that the formation of dense star clusters is a signature of efficient star formation, where $\\ge 20 \\%$ of the gas is converted into stars (e.g. \\cite{H80}, \\cite{L99}). Formation of numerous dense clusters can allow a galaxy to more efficiently convert gas into stars, thereby sustaining high SFRs in starburst systems."
    },
    "0206/astro-ph0206486_arXiv.txt": {
        "abstract": "{We report on coordinated thermal and optical measurements of trans-Neptunian object (20000) Varuna obtained in January-February 2002, respectively from the IRAM 30-m and IAA 1.5 m telescopes. The optical data show a lightcurve with a period of 3.176$\\pm$0.010 hr, a mean V magnitude of 20.37$\\pm$0.08 and a 0.42$\\pm$0.01 magnitude amplitude. They also tentatively indicate that the lightcurve is asymmetric and double-peaked. The thermal observations indicate a 1.12$\\pm$0.41 mJy flux, averaged over the object's rotation. Combining the two datasets, we infer that Varuna has a mean 1060$^{+180}_{-220}$ km  diameter and a mean 0.038$^{+0.022}_{-0.010}$ V geometric albedo, in general agreement with an earlier determination using the same technique. ",
        "introduction": "Our view of the outer Solar System has changed dramatically over the last decade, with the discovery of hundreds of objects beyond Neptune's orbit. These trans-Neptunian objects (TNOs), which can be classified in three dynamical groups from their orbital properties, are believed to have formed in the tenuous outskirts of the protoplanetary disk, and to have remained relatively unaltered since then (e.g. Jewitt \\& Luu 2000). As such, and since they are also thought to be the source of short-period comets (Duncan et al. 1988), these bodies are currently the subject of considerable interest. Because of their intrinsic faintness, however, the physical and chemical properties of TNOs are difficult to study. For most of them, physical observations are restricted to broad-band photometry, providing magnitudes, colors, and rotation periods (e.g. Doressoundiram et al. 2001, 2002). Only for a few of them have infrared spectra been acquired, with compositional diagnostics.  Among this population,  the ``classical TNO\" (20000)Varuna, discovered in November 2000 under the provisional designation of 2000 WR$_{106}$ (McMillan \\& Larsen 2000), has received special attention.  Thanks to prediscovery observations dating back to 1955 (Knoffel \\& Stoss 2000), its orbit is accurately known (to within 0.1 arcsec). With an apparent visual magnitude of about 20.3 (absolute visual magnitude, H= 3.7), it is one of the brightest known TNOs, being as of today surpassed only by (28978)Ixion (H=3.2). Visible photometry indicates that Varuna is moderately red (B-R $\\sim$ 1.5) (Jewitt \\& Sheppard 2002, hereafter JS02; Doressoundiram et al. 2002), and near-IR spectroscopy suggests the presence of water ice bands at 1.5 and 2.0 $\\mu$m (Licandro, Oliva \\& Di Martino, 2001). Jewitt, Aussel \\& Evans (2001, hereafter JAE01) reported the detection of Varuna at 850 $\\mu$m from JCMT observations, with a flux of 2.81$\\pm$0.85 mJy. From the combination of this thermal emission measurement with simultaneous optical observations, they inferred an equivalent circular diameter of 900$^{+129}_{-145}$ km and a red geometric albedo of p$_r$~=~0.070$^{+0.030}_{-0.017}$.  Shortly after, Farnham (2001) reported that Varuna's exhibits a rotational lightcurve, with a 0.5 mag amplitude and a single-peaked period of 3.17 hour, although periods of 2.78 and 3.67 hours could not be ruled out. The rotational behaviour of Varuna was extensively investigated by JS02 who reported a two-peaked R lightcurve with period 6.3442$\\pm$0.0002 hour and 0.42$\\pm$0.02 mag amplitude, and no rotational variations in the visible colors (B-V, V-R, V-I). They concluded that Varuna is probably an elongated, prolate body with a (projected) axis ratio as high as 1.5:1.  We present here additional combined observations of Varuna in the thermal and visible range, performed from two telescopes (IRAM-30m and IAA 1.5m, respectively), located on Pico Veleta, Sierra Nevada (southern Spain). The prime goal was to obtain an independent measurement of Varuna's thermal flux to confirm the single detection of JAE01.  A secondary, more difficult, objective was to search for rotational variability in the thermal flux. Indeed, a positive correlation of the thermal lightcurve with the optical lightcurve would indicate a shape effect (as is possibly the case for Ceres (Altenhoff et al. 1996) and Vesta (Redman et al. 1992)), while anticorrelation is the manifestation of albedo markings (an example is Pluto, Lellouch et al. 2000a). Unfortunately, our thermal observations did not prove of sufficient quality for this goal. ",
        "conclusions": "We have performed coordinated optical and thermal observations of Trans-Neptunian object (20000)Varuna. The optical data, acquired at the IAA 1.5 m telescope, show a clear lightcurve with a single-peaked period of 3.176$\\pm$0.010 hr, a mean V magnitude of 20.37$\\pm$0.08 and a 0.42$\\pm$0.01 magnitude amplitude. Phasing our observations with those of Jewitt \\& Sheppard (2002), we find a best fit period of 3.1788$\\pm$0.0001 hr. Our observations tentatively confirm an asymmetry in the lightcurve, as first reported by Jewitt \\& Sheppard. This would favor the hypothesis that the lightcurve is actually double-peaked with a 6.3576$\\pm$0.0002 hr period and predominantly due to an elongated shape of the object. The thermal data, obtained with the IRAM 30-m telescope, consist of five independent measurements of Varuna's 1.2 mm flux, sampling the optical lightcurve. These measurements are much too noisy to  distinguish a possible thermal lightcurve. Averaged together, they indicate a 1.12$\\pm$0.41 mJy flux at 1.2 mm, i.e. a 2.7 $\\sigma$ detection that adds to the 3.3 $\\sigma$ detection of Jewitt, Aussel \\& Evans (2001) at 0.8 mm and confirms the difficulty of this kind of observations. Assuming emissivity and thermophysical surface properties similar to Pluto's, the thermal data indicate a mean equivalent circular diameter of 1060$^{+180}_{-220}$ km. The associated albedos in the visible and the red are $p_{\\rm v}$~=~0.038$^{+0.022}_{-0.010}$ and $p_{\\rm r}$ = 0.049$^{+0.029}_{-0.013}$, respectively, consistent with the determination by JAE01. Taken together with the albedo measurement of 1993 SC ($p_{\\rm v}$~=~0.022$^{+0.013}_{-0.006}$) and the possible detection of 1996 TL$_{66}$ by Thomas et al. (2000), this indicates that the canonical 0.04 albedo adopted for size distribution studies is not invalid at this point. \\\\  {\\em Acknowledgments:} This research is partially based on data taken at the 1.5m telescope of Sierra Nevada Observatory which is operated by the Consejo Superior de Investigaciones Cientificas through the Instituto de Astrofisica de Andalucia. N.P. acknowledges funding from the FCT, Portugal (ref: SFRH/BD/1094/2000)."
    },
    "0206/astro-ph0206492_arXiv.txt": {
        "abstract": "Spatially resolved {\\it ROSAT} X-ray and ground-based optical data for the southwestern region of the Cygnus Loop SNR reveal in unprecedented detail the very early stages of a blast wave interaction with an isolated interstellar cloud. Numerous internal cloud shock fronts near the upstream flow and along the cloud edges are visible optically as sharp filaments of enhanced H$\\alpha$ emission. Faint X-ray emission is seen along a line of Balmer-dominated shock filaments north and south of the cloud with an estimated X-ray gas temperature of $1.2 \\times 10^{6}$ K (0.11 keV) corresponding to a shock velocity of 290 km s$^{-1}$. The main cloud body itself exhibits little or no X-ray flux. Instead, X-ray emission is confined along the northern and southernmost cloud edges, with the emission brightest in the downstream regions farthest from the shock front's current position. We estimate an interaction age of $\\sim$ 1200 yr based on the observed shock/cloud morphology.  Overall, the optical and X-ray properties of this shocked ISM cloud show many of the principal features predicted for a young SNR shock -- ISM cloud interaction. In particular, one sees shocklet formation and diffraction inside the inhomogenous cloud along with partial main blast wave engulfment. However, several significant differences from model predictions are also present including no evidence for turbulence along cloud edges, diffuse rather than filamentary [\\ion{O}{3}] emission within the main body of the cloud, unusually strong downstream [\\ion{S}{2}] emission in the postshock cloud regions, and confinement of X-ray emission to the cloud's outer boundaries. ",
        "introduction": "Supernova remnants (SNRs) shape and enrich the chemical and dynamical structure of the interstellar medium (ISM) which, in turn, affect the evolution of a SNR. Knowledge of just how SN generated shock waves travel through and interact with the ISM and interstellar clouds is fundamental to our understanding of the emission and dynamical details of this process.  Because of its large angular size (2.8$^{\\circ} \\times$ 3.5$^{\\circ}$), low foreground extinction ($E[B-V] = 0.08$ mag; \\citealt{parker67,fesen82}), and wide range of shock conditions, the Cygnus Loop is one of the best laboratories for studying the ISM shock physics of middle-aged remnants. At a distance of 440$_{-110}^{+150}$ pc \\citep{blair99}, it has a physical size of 21 $\\times$ 27 pc. Located 8.5$^{\\circ}$ below the galactic plane, the Cygnus Loop lies in a multi-phase medium containing large ISM clouds with a hydrogen density of $n=5-10$ cm$^{-3}$, surrounded by a lower density intercloud component of $n\\approx0.1-0.2$ cm$^{-3}$ \\citep{denoyer75}.  The currently accepted view of the Cygnus Loop is that it represents an ISM cavity explosion by a fairly massive progenitor star \\citep{mccray79,charles85,levenson99}. The cavity is presumably the result of strong stellar winds emanating from the high-mass progenitor. In this picture, the supernova shock has been traveling relatively unimpeded for a distance $\\approx$ 10 pc and has only relatively recently begun to reach the cavity walls. The interaction of the shock with the cavity walls is responsible for the remnant's observed radio, optical, and X-ray emission.  Previous studies of the Cygnus Loop have examined selected regions in the UV/optical and X-ray \\citep{ku84,hester86b,graham95,levenson96,danforth00}. These have shown that there are two distinctly different types of optical line-emission filaments present. The Cygnus Loop's brighter filaments are the result of shocked and subsequently radiatively cooled interstellar clouds whose preshock densities are many times that of the intercloud regions. Along with hydrogen and helium recombination line emissions, these filaments exhibit strong forbidden line emissions from oxygen, nitrogen, and sulfur, and are located downstream from the advancing shock front in postshock gas with temperatures $\\sim$ 10$^{5}$ K \\citep{fesen85}. The degree of postshock cooling (``incompleteness'') can strongly affect the relative strength of the line emissions, particularly the observed [\\ion{O}{3}] $\\lambda\\lambda$5007,4959 vs. H$\\beta$ emissions. In the case of the Cygnus Loop, like most other evolved SNRs, comparisons with model calculations show its bright filaments have shock velocities $\\approx$100 km s$^{-1}$.  Fainter, so-called Balmer-dominated filaments result when a high-velocity shock encounters partially neutral gas \\citep{chevalier78,chevalier80}. The collisionless shock accelerates and heats interstellar ions and electrons through electromagnetic plasma instabilities, while leaving neutral atoms unaffected. Subsequently, the neutral atoms, particularly neutral hydrogen, can be collisionally excited as well as collisionally ionized thus permitting the emission of Balmer photons with a narrow line profile width corresponding to the preshock gas temperature of T $\\sim 5000 - 10000$ K. However, neutral hydrogen in the postshock region can also undergo charge transfer thereby acquiring thermal energy and flow velocities similar to those of the shocked ions. Consequently, charge transfer to the shock-heated protons produces fast moving hydrogen atoms, which will emit Balmer photons with broad line profiles upon collisional excitation.  Other elements are also collisionally ionized and may also emit line photons. However, in the case of neutral atoms and relatively low-ionization ions, a line's luminosity is proportional to its ionization time, collisional excitation rate, and elemental abundance. This leads to relatively weak metallic lines compared to the hydrogen Balmer lines and thus Balmer-dominated filaments.  In general, lower density intercloud regions of the remnant experience higher velocity shocks and correspondingly show higher postshock temperatures. These intercloud regions are thus responsible for a remnant's X-ray and coronal line emissions \\citep{mckee75,ku84,teske90}. X-ray analyzes provide information on elemental abundances, shock front position, grain destruction, and other properties of the postshock gas. Furthermore, in cases where the postshock gas is fully ionized, X-ray emission can yield a direct measure of the shock front velocity.  Attempts to fit shock models to the observed optical line emission seen throughout the Cygnus Loop have sometimes been hindered by surprisingly large [\\ion{O}{3}]/H$\\beta$ ratios leading to the notion of incomplete shock emission \\citep{blair91}. Using IUE observations, \\citet{raymond80} found that much of the hydrogen recombination zone predicted by steady-flow models is absent, implying that the interaction is fairly young, with an incomplete postshock cooling and recombination zone. Similarly, small portions of NE limb Balmer-dominated filaments have been found to exhibit incomplete postshock cooling zones, apparently marking locations of increased ISM density and hence somewhat shorter cooling times.  Some of the results from these analyses are likely affected by other factors including limb projection effects, uncertain location of the associated forward shock, and the superposition of multiple shock fronts along the line of sight. Ideally, to compare observations with model emission calculations, one would like to observe a a single, isolated ISM cloud largely free of such complicating effects.  Towards this goal, \\citet{graham95} combined X-ray and optical data to study a small cloud in the southeast \\citep{fesen92} seen in the early stages of shock interaction. They found that the Balmer dominated emission, together with X-ray emission, traced out the shock front as it wrapped around the cloud. Their analysis together with optical images taken with the Hubble Space Telescope \\citep{levenson01} led them to picture the cloud, initially identified as small and isolated, as in fact an extension of a much larger cloud, which the blast wave is just now interacting with. A complex morphology of interacting shock fronts is seen where sharp filaments mark regions where the shock front is viewed edge-on and diffuse emission where the view is face-on.  Here we present optical and X-ray data on a small, isolated and recently shocked cloud located along the southwestern limb of the remnant. The shock-cloud interaction is viewed nearly edge-on, with the shock front visible both within and around the cloud. Using ground-based optical and ROSAT X-ray images and spectra, we present an analysis of the very early stages of this shock-cloud interaction. Our optical and X-ray observations are described in \\S\\ref{sect:observ}. We discuss the results in \\S\\ref{sect:res} and compare these properties to those of other previously studied regions of the Cygnus Loop in \\S\\ref{sect:disc}. In \\S\\ref{sect:conc}, we summarize our results and discuss its implications for the modeling os shock-cloud dynamics. ",
        "conclusions": "\\label{sect:conc}  {\\it ROSAT} X-ray data and ground based optical data for the southwestern region of the Cygnus Loop SNR show the early stages of the interaction of a blast wave with a cloud in unprecedented detail. From our study of this shocked cloud, we conclude the following:  1) The cloud began interacting with the shock $\\sim 1200$ yr ago. This is supported by the lack of a standing bow shock behind the cloud, the lack of a shock reconvergence point west of the cloud, and no evidence for instability formation along the edges of the cloud.  2) The optical morphology of the cloud is substantially different than what is seen in the brighter regions of the remnant. Whereas many of the brighter regions of the Cygnus Loop are the result of $\\sim$ 400 km s$^{-1}$ shocks hitting relatively higher density material, the low density and low shock velocity nature of this region stretches the postshock cooling zone resulting in the diffuse [\\ion{O}{3}] and clumpy [\\ion{S}{2}] emissions observed.  3) The cloud's X-ray emission structure is also unlike that seen in the brighter optical and X-ray regions of the remnant. Little or no X-ray emission is associated with the cloud itself, but there is bright X-ray emission associated with the northern and southern peripheries of the cloud. Furthermore, we derive a shock velocity of 290 km s$^{-1}$ which is significantly lower than shock velocities found in other parts of the Cygnus Loop remnant.  4) Small scale density fluctuations were found to exist within this ISM cloud which significantly altered the progression of the shock through the cloud. This is seen by the presence of multiple small scale shocks which are seen throughout the cloud.  In summary, this relatively isolated, low density cloud in the southwest limb of the Cygnus Loop has provided a revealing snapshot of the very early stages of a shock-cloud interaction. It shows how a cloud's initial density structure can strongly influence the observed optical and kinematic morphology of the postshock gas. Further analyses of some of the exquisite details of this shock-cloud interaction may provide additional diagnostics with regards to two- and three-dimensional models of shocks overrunning ISM clouds."
    },
    "0206/astro-ph0206037_arXiv.txt": {
        "abstract": "Observations indicate that massive stars in the Galaxy form in regions of very high surface density, $\\Sigma\\sim 1$ g cm\\ee.  Clusters containing massive stars and globular clusters have a column density comparable to this. The total pressure in clouds of such a column density is $P/k\\sim 10^8-10^9$ K cm\\eee, far greater than that in the diffuse interstellar medium or the average in giant molecular clouds. Observations show that massive star-forming regions are supersonically turbulent, and we show that the molecular cores out of which individual massive stars form are as well. The protostellar accretion rate in such a core is approximately equal to the instantaneous mass of the star divided by the free-fall time of the gas that is accreting onto the star (Stahler, Shu, \\& Taam 1980). The star-formation time in this {\\it Turbulent Core} model for massive star formation is several times the mean free-fall time of the core out of which the star forms, but is about equal to that of the region in which the core is embedded.  The high densities in regions of massive star formation lead to typical time scales for the formation of a massive star of about $10^5$~yr.  The corresponding accretion rate is high enough to overcome the radiation pressure due to the luminosity of the star.  For the typical case we consider, in which the cores out of which the stars form have a density structure $\\rho\\propto r^{-1.5}$, the protostellar accretion rate grows with time as $\\dot m_*\\propto t$.  We present a new calculation of the evolution of the radius of a protostar and determine the protostellar accretion luminosity.  At the high accretion rates that are typical in regions of massive star formation, protostars join the main sequence at about $20 M_\\odot$. We apply these results to predict the properties of protostars thought to be powering several observed hot molecular cores, including the Orion hot core and W3($\\rm H_2O$). In the Appendixes, we discuss the pressure in molecular clouds and we argue that ``logatropic'' models for molecular clouds are incompatible with observation. ",
        "introduction": "Massive stars are fundamental in the evolution of galaxies since they produce the heavy elements, energize the interstellar medium, and possibly regulate the rate of star formation.  Remarkably little is known about how massive stars form, however: the problem is difficult observationally because massive star formation occurs in distant, highly obscured regions, and it is difficult theoretically because of the many processes that must be included.  Even such a basic parameter as the time it takes to form a massive star has been uncertain.  This time scale, or equivalently, the protostellar accretion rate, affects the luminosity of the protostar (particularly for masses $m_*\\la 10 M_\\odot$ (Palla \\& Stahler 1992) and the strength of protostellar outflows (Richer et al. 2000).  Arguments based on extrapolating from low-mass star formation lead to formation times $\\tsf>10^6$ yr, a significant fraction of the main-sequence lifetime of the star (Bernasconi \\& Maeder 1996; McLaughlin \\& Pudritz 1997, hereafter MP97; Stahler et al. 2000).  Comparison with observations of hot molecular cores (Osorio, Lizano \\& D'Alessio 1999; Nakano et al. 2000) suggest substantially smaller time scales, $\\tsf\\la 10^5$ yr.  An analysis based on observations of protostellar outflows suggests $\\tsf\\sim 3\\times 10^5$ yr (Behrend \\& Maeder 2001). The small ($\\sim 1\\times 10^{6}\\:{\\rm yr}$) spread in ages of stars in the Orion Nebula Cluster (Palla \\& Stahler 1999), where there is no evidence that the higher mass stars have formed systematically later compared to the lower-mass population, sets an upper limit of $\\tsf\\la 1\\:{\\rm Myr}$ in this case.  What has been lacking is an adequate understanding of how the formation time is governed by the conditions in the gas out of which the star forms.   Our understanding of low-mass star formation is on a far better footing, since it has received much more observational and theoretical attention (Shu, Adams \\& Lizano 1987).  Low-mass stars form by accreting gas from a molecular ``core'' in which gravity overcomes thermal and nonthermal (magnetic and turbulent) pressure gradients. Shu (1977) considered the collapse of a singular isothermal sphere, finding \\begin{equation} \\label{eq:shu} \\dot{m}_*=0.975\\frac{\\cth^3}{G}=4.36\\times 10^{-6} \\left(\\frac{T}{20\\:{\\rm K}}\\right)^{3/2}\\smyr, \\end{equation} where $\\cth$ is the isothermal sound speed and the numerical evaluation assumes $n_{\\rm He}=0.2 n_{{\\rm H_2}}$. Observed temperatures of $10$ to $20\\:{\\rm K}$ in regions of low-mass star formation imply accretion rates of about $10^{-6}$ to $10^{-5}\\:{M_\\odot\\:yr^{-1}}$, consistent with the inferred values of $\\tsf$ for low-mass stars in these regions (Lada 1999).  There are two difficulties in extending this theory to high-mass stars.  The first, discussed in some detail by Stahler et al. (2000), is that the predicted accretion rate depends only on the temperature of the gas. Once massive stars form, the gas may be heated to temperatures $\\sim 50$ to $100\\:{\\rm K}$, but the first massive stars that form in a region will emerge from gas at $10-20$~K and will have low accretion rates and correspondingly long formation times. The second difficulty is feedback from the massive stars. Since the Kelvin-Helmholtz contraction time is less than the accretion time for massive stars, they evolve along the main sequence while accreting.  Massive protostars are thus very luminous, and it has been suggested that the radiation pressure and ionization they produce can halt the accretion and determine the upper limit of the stellar mass function (Larson \\& Starrfield 1971; Kahn 1974; Wolfire \\& Cassinelli 1987; Jijina \\& Adams 1996). This feedback is so strong that it is impossible to form stars as massive as those observed if the accretion is assumed to be spherical, and the discrepancy grows as the accretion rate is reduced.  These considerations have motivated the radical suggestion that massive stars form via the coalescence of low-mass stars in order to achieve a more rapid build up of the final stellar mass (Bonnell, Bate \\& Zinnecker 1998).  Recently, McKee \\& Tan (2002; hereafter MT) addressed the accretion-rate problem.  The first step in resolving the problem of the apparently low accretion rates is to realize that, although equation (\\ref{eq:shu}) was derived for an isothermal gas, it should hold approximately when nonthermal support due to magnetic fields and turbulence is included as well (Stahler, Shu \\& Taam 1980; Shu et al. 1987). Observed cores have turbulent motions that increase systematically with radius (Larson 1981; Caselli \\& Myers 1995), and this leads to an increase in the accretion rate with time (Myers \\& Fuller 1992; Caselli \\& Myers 1995; MP97). Larger signal speeds allow for the hydrostatic support of denser gas cores, which then have shorter free-fall times and thus greater accretion rates once they become unstable. We term this the {\\it turbulent core} model for massive star formation.  The second step in resolving the accretion rate problem is the recognition that massive stars form in regions of very high pressure and density.  MT showed that for typical pressures in regions of massive star formation (Plume et al. 1997), stars form in a time of order $10^5$~yr.  This result for the time scale is somewhat longer than that of Osorio et al. (1999), who inferred stellar masses and accretion rates by comparing calculated spectra with observations.  The purpose of this paper is several fold. First, we determine the relation between the pressure in a molecular cloud and its surface density $\\Sigma$. We show that observed regions of massive star formation, both Galactic and extragalactic, have $\\Sigma\\sim 1$ g cm\\ee, corresponding to mean pressures $\\bar P/k\\sim 4\\times 10^8$ K cm\\eee. Second, we extend the self-similar theory presented by MT to allow for magnetic fields and for a thermally supported core.  Finally, we determine the radius and luminosity of accreting protostars, and use observed hot core luminosities to predict accretion rates and masses of several nearby massive protostars. ",
        "conclusions": "We have developed a model (that we term the Turbulent Core Model) for the formation of massive stars, which is an extension of the classic paradigm of low-mass star formation (Shu, Adams, \\& Lizano 1987) and is to be contrasted with models involving competitive accretion (Bonnell et al. 1997, 2001) or stellar collisions (Bonnell et al. 1998). The principal motivations for the latter models are the short formation timescales (and correspondingly high accretion rates) mandated by observations of short star cluster formation times (Palla \\& Stahler 1999) and theoretical considerations of radiation pressure feedback (Wolfire \\& Cassinelli 1987). Such accretion rates are difficult to justify in the standard picture of isothermal core collapse (Shu 1977). Collapse from cores with nonthermal pressure support can involve faster accretion rates (Stahler et al. 1980), but there has been no self-consistent theory for predicting both the normalization of the expected accretion rates and their evolution. This has led to a vast range of massive star accretion rates ($10^{-6}-10^{-2}\\smyr$) being considered in the recent literature (e.g. Bernasconi \\& Maeder 1996; MP97; Nakano et al. 2000).  Conventional theories of massive star formation face further problems: massive stars form preferentially in the centers of stellar clusters (Bonnell \\& Davies 1998) where the crowded environment makes it difficult to understand the existence of massive pre-stellar cores, and the high densities and pressures lead to a small thermal Jeans mass that is only a fraction of a solar mass (Bonnell et al. 1998).  The above criticisms, taken together with observational hints that massive stars may form differently from low-mass stars (central concentration in star clusters, high degree of equal-mass binarity, complex morphology and extreme energetics of outflows), have motivated collisional and competitive accretion models. However, these are also not without their problems: as noted by Bonnell et al. (1998), competitive (Bondi-Hoyle) accretion is suppressed for stellar masses above $\\sim 10\\sm$ because of radiation pressure feedback, while stellar collisions require extreme $\\gtrsim 10^{8}\\:{\\rm pc^{-3}}$ stellar densities, which have never been observed.  The Turbulent Core Model for massive star formation overcomes the difficulties of the standard accretion scenarios by incorporating the effects of the supersonic turbulence and high pressures observed in massive star-forming regions (Plume et al. 1997). The high pressures mean that cores that become unstable are necessarily very dense and small, leading to high accretion rates and no over-crowding. The turbulent and nonthermal nature of cores gives them substructure and thus the protostar's accretion rate will exhibit fluctuations about the mean. Note that while clumpiness in the core is an attribute of our model for massive star formation, it is not a {\\it requirement}, as in the collisional model of Stahler et al (2000). The turbulence and nonthermal support of the clump (protocluster) gas determines the IMF for stars above about a solar mass, whereas the Bonnor-Ebert mass is important in determining the IMF at lower masses. While our theory does not aim to predict the IMF, we note that the empirical core mass function is not too different from the stellar one (Testi \\& Sargent 1998; Motte et al. 2001), at least up to $\\sim 5\\sm$, the maximum mass probed by these observations. More massive cores, with relatively simple, centrally-concentrated morphologies have been observed, e.g. the Orion hot core (Wright et al. 1992) and W3($\\rm H_2O$) (Wyrowski et al. 1999). The stellar IMF would then be set by the core mass function, modulated by $\\ecore(m_*)$. In a model in which $\\ecore$ is set by the feedback from bipolar outflows, it is found to have a relatively weak dependence on stellar mass (Matzner \\& McKee 2000; Paper II).   A natural question is why gravitational fragmentation does not continue down to the thermal Jeans (sub-solar) mass scale inside a massive core, leading to the formation of a cluster of low-mass stars instead of a single massive star. In the self-similar model we have presented, clumps and cores exhibit density fluctuations on all scales down to the thermal Jeans mass. We have assumed that clumps and cores are relatively long-lived (over at least several dynamical times), which implies that most of the density fluctuations existing at a particular instant are gravitationally stable. How good is this assumption? The age spread of pre-main sequence stars in the Orion nebula cluster is of order 1~Myr, which is about an order of magnitude greater than the expected free-fall time of a typical Plume et al. (1997) clump. The approximately spherical morphologies of many clumps observed by Shirley et al. (2002) suggests that they have existed for at least a dynamical timescale. If the onset of gravitational instability is a relatively rare phenomenon in this environment (i.e., $\\la$ 20\\% of the mass of a clump is undergoing gravitational collapse at any time), as we have assumed, then it is perhaps not surprising that gravitational fragmentation in a collapsing core is unlikely. Gravitational fragmentation is further suppressed by the tidal field of the embedded stars and protostars.   The main results of this paper are the following:   1. The surface densities in observed regions of massive star formation and in star clusters that contain, or did contain, massive stars, are typically within a factor 4 of $\\scl\\sim 1$ g cm\\ee. The corresponding mean pressures are $\\bar P/k\\simeq G\\scl^2\\sim 10^8-10^9$~K~cm\\eee, much greater than in the diffuse ISM or the typical location in a GMC.   2. Cores have column densities similar to that of the clump in which they are embedded.  The mean density of a core significantly exceeds that of its natal clump, and the radius is less than the the tidal radius.  3. Cores that form massive stars are supersonically turbulent; there is no need for the gas to become subsonic in order for star formation to occur.  4. The star-formation time is several times the mean free-fall time of the core out of which the star forms, but is about equal to that of the region in which the core is embedded.  5. The time for a massive star to form in a typical region of of massive star formation is about $10^5$~yr, with a weak ($m_{*f}^{1/4}$) dependence on stellar mass and a somewhat stronger dependence on the surface density of the clump in which it is forming ($\\scl^{-3/4}$). This timescale is short compared to estimated cluster formation times, but long compared to the ages of observed supernova remnants, which sometimes have been invoked as star formation triggers.  6. The corresponding accretion rate, approaching $10^{-3}\\smyr$ for the most massive stars, is high enough to overcome the radiation pressure due to the luminosity of the star.  7. For the typical case we consider, in which the cores out of which the stars form have a density structure $\\rho\\propto r^{-1.5}$, the protostellar accretion rate grows with time as $\\dot m_*\\propto t$. These density structures are consistent with observed clumps and cores, while in an appendix we have shown that logatropic models are inconsistent.  8. The rate at which a core accretes mass from the ambient clump is comparable to the rate at which it processes matter into a star. Once the star has formed, subsequent Bondi-Hoyle accretion is negligible, particularly for massive stars.  9. Presenting a calculation of the evolution of the radius of a protostar, we determine the protostellar accretion luminosity. When the (eventually) massive protostar is still less than a few solar masses, this luminosity can be several hundred to a thousand solar luminosities. Massive protostars join the main sequence at around $20\\sm$.  10. Application to observations of the Orion hot core suggests a current protostellar mass of between about 11 and 18~$\\sm$ and an accretion rate of a few$\\times 10^{-4}\\smyr$. Similar properties are estimated for W3($\\rm H_2O$), the Turner-Welch object.  The incorporation of feedback, including protostellar outflows, ionization and radiation pressure, is the subject of Paper II. In particular the question of when feedback prevents accretion is addressed. The implications of this work for star cluster formation will be examined in a future paper (see Tan \\& McKee 2002b for an initial discussion)."
    },
    "0206/astro-ph0206201_arXiv.txt": {
        "abstract": "{ We imaged the error box of a gamma-ray burst of the short (0.5 s), hard type (GRB 000313), with the BOOTES-1 experiment in southern Spain, starting 4 min after the $\\gamma$--ray event, in the $I$-band. A bright optical transient (OT 000313) with $I$ = 9.4 $\\pm$ 0.1 was found in the BOOTES-1 image, close to the error box (3$\\sigma$) provided by BATSE. Late time $VRIK^\\prime$-band deep observations failed to reveal an underlying host galaxy. If the OT 000313 is related to the short, hard GRB 000313, this would be the first optical counterpart ever found for this kind of events (all counterparts to date have been found for bursts of the long, soft type). The fact that only prompt optical emission has been detected (but no afterglow emission at all, as supported by theoretical models) might explain why no optical counterparts have ever been found for short, hard GRBs. This fact suggests that most short bursts might occur in a low-density medium and favours the models that relate them to binary mergers in very low-density environments. ",
        "introduction": "Gamma Ray Bursts (GRBs hereafter) are flashes of cosmic high energy photons, and they remained for 25 years one of the most elusive mysteries for high energy astrophysicists, the main problem being the lack of knowledge about the distance scale. The detection of counterparts at other wavelengths for the long duration, soft GRBs, revealing their cosmological origin (see van Paradijs et al. 2000 for a recent review). Thus, counterparts to about 30 bursts have been discovered so far with about 25 redshifts measured, but all of them belong to the so called long duration ($\\sim$ 20 s), soft bursts class that comprises about 75\\% of all GRBs (Mazets et al. 1981). There are evidences that the two classes of bursts are different: whereas long bursts have softer spectra, short bursts have harder spectra (Dezalay et al. 1996). The latter ones comprise about 25\\% of all GRBs (Kouvelioutou et al. 1993) and their origin still remain a puzzle. No counterparts at longer wavelengths have been found yet in spite of intense efforts in order to detect the optical, infrared and radio counterparts to several short, hard bursts (Kehoe et al. 2001, Gorosabel et al. 2002, Hurley et al. 2002, Williams et al. 2002). Therefore, one of the remaining GRB mysteries is whether the origin of the two populations are substantially different from one another.  \\begin{figure}[t] \\begin{center} \\resizebox{8cm}{4.1cm}{\\includegraphics{000313_fig1.ps}} \\resizebox{8cm}{4.1cm}{\\includegraphics{000313_fig2.ps}} \\caption{$2.1^{\\circ} \\times 1.1^{\\circ}$ fields in the $I$-band containing a fraction of the BATSE GRB error box. 13 Mar, UT 21 h 17 min ({\\it upper panel}) and 13 Mar, UT 22 h 13 min ({\\it left panel}). The position of the OT 000313 is marked with an arrow at the center of the image. North is upward and East is to the left. The limiting magnitude is $I$ = 12.0 for the second image.} \\end{center} \\end{figure}   Here we present the results of a follow-up observation for one of these short/hard events. GRB 000313 was detected on 13 March 2000, UT 21 h 13 min 04 s by the Burst and Transient Source Experiment (BATSE) instrument aboard the Compton Gamma-ray Observatory (trigger number 8035). The  single--peaked gamma--ray burst lasted 0.768 $\\pm$ 0.458 s, showed substructure and was possibly detected below 50 keV. It reached a peak flux (50-300 keV) of $\\sim$ 1.2 ($\\pm$ 0.1) $\\times$ 10$^{-7}$ erg cm$^{-2}$ s$^{-1}$ and a fluence ($\\geq$ 25 keV) of $\\sim$ 2.6 ($\\pm$ 1.7) $\\times$ 10$^{-7}$ erg cm$^{-2}$. The gamma-rays properties as well as the duration of the burst make from GRB~000313 a clear short/hard gamma-ray burst. The original BATSE position reported through the GCN/BACODINE Network (Barthelmy et al. 1998) was R.A.(2000) = 13h 31m, Dec(2000) = +27$^{\\circ}$ 10$^{\\prime}$ which was later on refined to R.A.(2000) = 13h 11m, Dec(2000) = +10$^{\\circ}$ 14$^{\\prime}$. ",
        "conclusions": "\\label{conclusiones}  If the OT 000313 is indeed related to the short GRB 000313 then, the fact that only prompt optical emission has been observed (and no optical afterglow emission) might explain the fact that no other optical counterparts for the short GRB class has been detected, favouring the models that relate them to binary mergers in galactic haloes or in the intragalactic medium. Given the fact that the distance distribution of NS-NS mergers depends on the mass of the host galaxy (Fryer, Woosley and Hartmann 1999), deeper observations of the OT 000313 error box might help to better constraint the distance and/or the mass of the host galaxy."
    },
    "0206/nucl-th0206053_arXiv.txt": {
        "abstract": " ",
        "introduction": "The rapid advance in nuclear reactions using rare isotopes has opened up several new frontiers in nuclear sciences\\cite{tanihata,liudo,pak97,lkb98,sjy,betty,udo,ditoro}. In particular, the intermediate energy heavy rare isotopes currently available at the National Superconducting Cyclotron Laboratory (NSCL/MSU) and the more energetic ones to be available at the future Rare Isotope Accelerator (RIA) in the United States provide a unique opportunity to explore novel properties of dense neutron-rich matter that was not in reach in terrestrial laboratories before. This exploration will reveal crucial information about the equation of state (${\\rm EOS}$) of neutron-rich matter. To understand the latter and its astrophysical implications, such as, the origin of elements, structure of rare isotopes and properties of neutron stars are presently among the most important goals of nuclear sciences. The ${\\rm EOS}$ of neutron-rich matter of isospin asymmetry $\\delta\\equiv (\\rho_n-\\rho_p)/(\\rho_n+\\rho_p)$ can be written as \\begin{equation}\\label{ieos} e(\\rho,\\delta)= e(\\rho,0)+E_{sym}(\\rho)\\delta^2 \\end{equation} within the parabolic approximation (see e.g., \\cite{lom}), where $e(\\rho,0)$ is the energy per nucleon in isospin symmetric nuclear matter. To study the density dependence of nuclear symmetry energy $E_{sym}(\\rho)$ has been a longstanding goal of extensive research with various microscopic and/or phenomenological models over the last few decades, e.g., \\cite{bru67,siemens70,fri81,chin77,serot,har87,mut87,wiringa88,bom91,akm98,hub98,my98,lee98}, for a recent review, see, e.g. \\cite{lom}. All models are able to give the symmetry energy at normal nuclear matter density $E_{sym}(\\rho_0)$ of about $34\\pm 4$ MeV in agreement with that extracted from the atomic mass data. However, the predicted results on the density dependence, especially at high densities, are extremely diverse and often contradictory. At subnormal densities, all models predict that the symmetry energy increase with density although the growing rate is rather model dependent. To extract the density dependence of symmetry energy at low densities, several approaches have recently been proposed. These include measuring the thickness of neutron-skins of rare isotopes\\cite{oya} and heavy stable nuclei, such as $^{208}Pb$\\cite{brown,hor,typel,fur}, isospin fractionation\\cite{lik97,xu,tan} as well as nuclear collective flow\\cite{ibuu2,sca99,li00} in intermediate energy heavy-ion collisions. Above normal nuclear matter densities, even the trend of the density dependence of $E_{sym}(\\rho)$ is still controversial. Theoretical results can be roughly classified into two groups, i.e., a group where the $E_{sym}(\\rho)$ rises monotonously and one in which it falls with the increasing density above about twice the normal nuclear matter density. The predictions also depend on the nature and form of effective interactions used. For instance, within the same Hartree-Fock approach using about 25 Skyrme and Gony effective interactions that have been widely used successfully in studying saturation properties of symmetric nuclear matter and nuclear structures near the $\\beta$ stability valley, the calculated symmetry energies were found to fall approximately equally into the two groups\\cite{brown,mar}.  The density dependence of nuclear symmetry energy, especially at high densities\\cite{kut}, has many profound consequences for various studies in astrophysics\\cite{lat00,pra97,bom01}. In particular, an increasing $E_{sym}(\\rho)$ leads to a relatively more proton-rich neutron star whereas a decreasing one would make the neutron star a pure neutron matter at high densities. Consequently, the chemical composition and cooling mechanisms of protoneutron stars\\cite{lat91,sum94}, critical densities for Kaon condensations in dense stellar matter\\cite{lee96,kurt2}, mass-radius correlations\\cite{prak88,eng94} as well as the possibility of a mixed quark-hadron phase\\cite{kurt3} in the cores of neutron stars will all be rather different. The high density (HD) behaviour of nuclear symmetry energy $E_{sym}(\\rho)$ is very important for understanding many interesting astrophysical phenomena, but it also subjects to the worst uncertainty among all properties of dense nuclear matter\\cite{kut}. The fundamental cause of the extremely uncertain HD behaviour of $E_{sym}(\\rho)$ is the complete lack of terrestrial laboratory data to constrain directly the model predictions. We explore in this work the possibility of using high energy heavy-ion collisions to probe the HD behaviour of $E_{sym}(\\rho)$. A brief report of this work can be found in ref. \\cite{li02}. The $E_{sym}(\\rho)$ is found to affect significantly several aspects of high energy heavy-ion collisions. In particular, the neutron/proton ratio of HD nuclear matter formed in these collisions is determined mainly by the HD behaviour of $E_{sym}(\\rho)$. Several promising signatures of the HD behaviour of $E_{sym}(\\rho)$ are also investigated. Besides the {\\it quantitative} observables, we also study one special {\\it qualitative} signature of the HD behaviour of the symmetry energy, i.e., the local minimum in the excitation function of nuclear transverse flow. The paper is organized as follows. In Section 2, we discuss the equation of state (${\\rm EOS}$) for neutron-rich matter, in particular, the HD behaviour of nuclear symmetry and its effects on the neutron/proton ratio in neutron stars. In section 3, we study the neutron/proton ratio and its dependence on the symmetry energy in HD hadronic matter formed in high energy heavy-ion collisions within an isospin-dependent hadronic transport model. In section 4, we explore several experimental probes of the HD behaviour of nuclear symmetry energy. Finally, we summarize in section 5. ",
        "conclusions": "In summary, the HD behaviour of nuclear symmetry energy has been puzzling physicists for decades. In this work, high energy heavy-ion collisions are proposed as a novel means to solve this longstanding problem. Within an isospin dependent hadronic transport model using two representative density dependent symmetry energy functions predicted by many body theories, it is shown that the isospin asymmetry of HD nuclear matter formed in high energy heavy-ion collisions is uniquely determined by the HD behaviour of $E_{sym}(\\rho)$. It is demonstrated that the isospin separation instability indeed can happen in high energy heavy-ion collisions. We explored several promising experimental probes for the HD symmetry energy. Among them, four quantitatively sensitive observables, the $\\pi^-/\\pi^+$ ratio, nuclear transverse collective flow and its excitation function as well as the neutron-proton differential collective flow are studied. A precursor of the possible isospin separation instability in dense neutron-rich matter is predicted to appear as the local minima in the excitation functions of the flow parameter for both neutrons and protons above the pion production threshold. This precursor can be used as a unique signature of the isospin dependence of the nuclear {\\rm EOS}. This {\\it qualitative} signature is advantageous over the other {\\it quantitative} observables because the latter are often affected by several ingredients of the reaction dynamics and experimental uncertainties. It should be noted that the present study is based on a momentum-independent transport model. The momentum-dependent potential is known to be important for a reliable description of global momentum distributions. Nevertheless, we expect that all qualitative features, such as the existence of the minimum in the excitation function of transerse flow, are not affected by the momentum-dependence. However, the quantitative results, such as the exact location of the flow minimum, are affected by the momentum-dependence of the nuclear potential. A more refined study using a momentum-dependent isoscalar potential is underway and the results will be reported elsewhere. It should also be mentioned that the fragmentation mechanism of isospin asymmetric nuclear matter is still not clear and remains a hot topic of current studies\\cite{li02a}. The different fragmentation mechanisms might affect the neutron/proton ratio of the observed free nucleons, and thus the neutron-proton differential flow. It would thus be useful to measure also observables associated with fragments together with the single particle observables examined in this study. It still remains a theoretical challenge to develop a fully quantum transport theory that deals with both particle production and cluster formation in a numerically tractable manner. The present work shows the need and promise of working in this direction. Experimental measurements of the proposed observables will provide the first terrestrial data to constrain stringently the HD behaviour of nuclear symmetry energy. Future comparisons between the experimental data and model calculations will allow us to extract crucial information about the {\\rm EOS} of dense neutron-rich matter.  This work was supported in part by the National Science Foundation Grant No. PHY-0088934 and Arkansas Science and Technology Authority Grant No. 00-B-14."
    },
    "0206/astro-ph0206084_arXiv.txt": {
        "abstract": "{We present optical $R$-band light curves of the gravitationally lensed quasar \\object{SBS1520+530} derived from data obtained at the Nordic Optical Telescope.  A time delay of $130\\pm3$ days (1$\\sigma$) is determined from the light curves. In addition, spectra of \\object{SBS1520+530} obtained at the Keck Observatory are spatially deconvolved in order to extract the spectrum of the faint lensing galaxy, free of any contamination by the light from the bright quasar images. This spectrum indicates a lens redshift $z=0.717$, in agreement with one of the absorption systems found in the quasar spectra. The best mass model of the system includes a second nearby galaxy and a cluster of galaxies in addition to the main lensing galaxy. Adopting this model and an $\\Omega=0.3$, $\\Lambda=0.7$ cosmology, our time-delay measurement yields a Hubble constant of H$_{0}=51 \\pm 9\\, {\\rm km}~{\\rm s^{-1}}~{\\rm Mpc^{-1}}$ (1$\\sigma$ error). ",
        "introduction": "Prompted by the successful optical measurements of time delays in PG~1115+080 (Schechter et al. \\cite{Sche}), a photometric monitoring campaign of gravitationally lensed quasars has been carried out at the Nordic Optical Telescope (NOT) with the aim of measuring time delays for many gravitationally lensed quasars.  This observable is of crucial importance to study the mass distribution in lensing galaxies and to determine $H_{0}$ using the method proposed by Refsdal (1964).  Two time  delays have  been measured so  far at  the NOT: one  for the doubly imaged quasar B~1600+434, lensed by an edge-on spiral (Burud et al. \\cite{Burud}) and another for the two (summed) components RX~J0911+0550 (Hjorth et al. \\cite{Hjorth}).  In  this paper we  will present our third  time delay measurement: that of  \\object{SBS1520+530}, a  doubly imaged  BAL quasar  at   $z=1.86$.   \\sbs\\,   was  discovered  by   Chavushyan  et al. (\\cite{Chavushyan}) as a  double quasar with angular separation of 1.56\\arcsec.   The   lensing  galaxy  was  detected   by  Crampton  et al.(\\cite{Crampton}) on $H$-band images obtained by using the adaptive optics  system  of the  Canada-France-Hawaii  Telescope.  The  lensing galaxy is also well resolved on optical images obtained at the NOT and deconvolved using the MCS  deconvolution algorithm (Magain, Courbin \\& Sohy \\cite{Magain}), and on public Hubble Space Telescope near-IR data (Faure  et al.  \\cite{Faure}).  The redshift  of  the galaxy  however, remains  unknown. With  the aim  of measuring  this redshift,  we have obtained a spectrum of \\sbs\\, with ESI at the Keck observatory.  The observations  and data  reduction of the  images are  presented in Sect.~\\ref{sect:data} and \\ref{sect:phot}. The  time delay measurement is described in Sect.~\\ref{sect:timedelay} while the analysis  of the spectroscopic data  is explained  in  Sect.~\\ref{sect:sbsspec}. This  section  also includes  a discussion  on the  spectral differences  between  the two quasar images.   Mass  models  and  the estimate  of  the  Hubble constant    are    discussed    in    Sect.~\\ref{sect:h0}.     Finally, Sect.~\\ref{sect:discussion} summarises the main results. ",
        "conclusions": "\\label{sect:discussion}  The  time delay  $\\Delta t  =  130 \\pm3$ ~days (1  $\\sigma$) has  been measured for the first time, in  the lensed quasar \\sbs\\, on the basis of $R$-band  images obtained with  the NOT.  Keck spectroscopy  of the lensing  galaxy  strongly  suggests  that  the  absorption  system  at $z=0.717$  is  associated  with  the  lensing  galaxy.  Applying the detailed mass model presented by Faure et al. (\\cite{Faure}) we derive a mean value H$_0= 51 \\pm9~ \\rm km s^{-1}Mpc^{-1}$. When only the main lensing galaxy is used in the mass model H$_0$ increases to 63 $\\pm 9~ \\rm kms^{-1}Mpc{-1}$.  The fit of the wo different models depends on the flux ratio between the two images, the model with only one galaxy only gives a good fit when the high flux ratio measured from the emission lines is applied. Determining precisely the flux ratio is therefore crucial to further improve the modeling of the system and hence the precision on H$_0$. Finally, a possible systematic error adds to the uncertainty in H$_0$ if the central mass concentration is not isothermal.  External  variations,  probably   due  to  microlensing  effects,  are observed    on     the    time    delay     shifted    light    curves (Fig.~\\ref{lightcurves}).   Part of this  effect can  be modeled  as a linear  term,  or  corrected  for  with the  iterative  algorithm  for measuring  time delays.  However,  significant external  variations of time scales of $\\sim$50 days do remain once these corrections are made.  Microlensing is important in \\sbs\\, as suggested both by the high frequency variations in the light curves and by the spectra of the quasar images. With light curves in only one band we can not efficiently disentangle between microlensing and intrinsic variations of the quasar. Such a work shall be possible using the colour information provide by a monitoring of the object in spectroscopy. The issue in conducting such a spectrophotometric monitoring for \\sbs\\, is double: (1) to measure the true flux ratio between the quasar images and discriminate between lens models and (2) to use microlensing to infer constraints on the distribution of micro-lenses in galaxy (L) and/or to reconstruct the energy profile of the central AGN in the source."
    },
    "0206/astro-ph0206421_arXiv.txt": {
        "abstract": "We investigate the relationship between the dark matter and baryons in the linear regime. This relation is quantified by the so-called ``filtering scale''. We show that a simple gaussian ansatz which uses the filtering scale provides a good approximation to the exact solution. ",
        "introduction": "If humans had dark matter vision, cosmology would be a solved problem by now. However, this is not the case, and in our attempt to understand the distribution of matter in the universe we need to rely on baryons - stars and gas - that trace the underlying distribution of the dark matter.  The baryons, though, are not the perfect tracer, since they are subject to other forces in addition to gravity. In this paper we restrict our attention to the linear regime, but even in this simplest physical situation the relationship between the dark matter and baryons (i.e.\\ cosmic gas) is nontrivial, because gas pressure will erase fluctuations on small scales.  The effect of gas pressure on small fluctuations is canonically described by the Jeans scale, which is defined as the scale at which the gravity force equals the gas pressure force. On large scales gravity wins, and small fluctuations grow exponentially, while on small scales, gas pressure turns all fluctuations into sound waves.  However, in the expanding universe the Jeans scale becomes essentially irrelevant, because gravitational instability leads only to slow, power-law growth of fluctuations. Let us consider the following simple thought experiment: in a universe with linear dark matter fluctuations, the gas is instantaneously heated to high temperature. The Jeans scale also increases by a large factor instantaneously, whereas it takes about a Hubble time for the fluctuations in the gas to respond to the changed pressure force. Thus, the instantaneous value of the Jeans scale does not correspond to the characteristic scale over which the fluctuations are suppressed, but instead the suppression scale - which we call the ``filtering'' scale following Gnedin \\& Hui (1998) - depends on the whole previous thermal history. Only in the unphysical case of temperature evolving as an exact power-law of the scale factor at all times does the Jeans scale become proportional to the filtering scale (Bi, Borner, \\& Chu 1992; Fang et al.\\ 1993). However, incorrect expressions for the pressure force filtering have been used even until quite recently (c.f.\\ Choudhury, Padmanabhan, \\& Srianand 2001)   In this paper we investigate the role of the filtering scale further. Specifically, it has been suggested (Gnedin \\& Hui 1998; Gnedin 1998) that the relationship between the gas density fluctuation $\\delta_B$ and the dark matter fluctuation $\\delta_X$ in the Fourier domain can be approximated by the following simple expression: \\begin{equation} {\\delta_B\\over\\delta_X} = e^{\\displaystyle-k^2/k_F^2}, \\label{app} \\end{equation} where $k$ is a wavenumber, and $k_F$ is the wavenumber corresponding to the filtering scale. Our goal will be to investigate the range of validity of this approximation. ",
        "conclusions": "We have shown that the gaussian approximation (\\ref{app}) provides a reasonably good fit (better than 10\\%) to the rms density fluctuations in the gas, and for the shape of the gas power spectrum for $k<0.9k_F$. It is important to underscore that there is no known physical reason why this approximation works so well, so it should be considered as a mathematical coincidence. Notwithstanding, the gaussian approximation can be used in semianalytical models of the Lyman-alpha forest and early universe, when a full solution to equations (\\ref{let}) is not practical."
    },
    "0206/astro-ph0206167.txt": {
        "abstract": "We have observed with the 30-m IRAM telescope, the CSO telescope and the ISO\\footnotemark[4] satellite ({\\em Infrared Space Observatory}) the rotational lines of CO at millimeter, submillimeter and far infrared wavelengths in the direction of C-rich stellar objects at different stages of evolution :  CRL 2688 (a very young Proto-Planetary Nebula), CRL 618 (a Proto-Planetary Nebula), and NGC 7027 (a young Planetary Nebula). Several changes in the longwave emission of CO and other molecules are discussed here in relation with the degree of evolution of the objects. In the early stages, represented by CRL 2688, the longwave emission is dominated by CO lines. In the intermediate stage, CRL 618, very fast outflows are present which, together with the strong UV field from the central star, dissociate CO. The released atomic oxygen is seen via its atomic lines, and allows the formation of new O-bearing species, such as \\water and OH. The abundance of HNC is enhanced with respect to HCN as a result of the chemical processes occurring in the photo-dissociation region (PDR). At this stage, CO lines and [O$\\small I$] lines are the dominant coolants, while the cooling effect of [C$\\small{II}$] is rising. At the Planetary Nebula stage, NGC 7027, large parts of the {\\em old} CO AGB material have been reprocessed. The spectrum is then dominated by atomic and ionic lines. New species such as CH$^{+}$ appear. Water has probably been reprocessed in OH. ",
        "introduction": "Solar type stars remain in the asymptotic giant branch (AGB) for about $\\sim 10^{6}-10^{7}$ years, losing mass through stellar winds ($\\dot{M} \\sim 10^{-4}-10^{-7}$ M$_{\\odot}/$yr, ejection velocity $\\sim 5-25$ \\kms, see, e.g., Loup \\etal 1993). During this phase the ejected material progressively forms an expanding circumstellar envelope (CSE). Molecular species are easily formed in the innermost regions of the CSE. Some of these molecules aggregate onto  dust grains which under the action of radiation pressure accelerate and push the remaining gas producing an expanding dusty molecular envelope. UV photons from the interstellar radiation field dissociate gas-phase molecules in the external layers of the envelopes allowing new chemical reactions and the production of new molecular species. These chemical changes have been very well studied through observations and models in the case of IRC+10216, the prototypical AGB carbon-rich star (Glassgold 1996; Cernicharo, Gu\\'elin \\& Kahane 2000 and references therein). Additional changes in the composition of the CSE do occur when the central object starts its evolution towards the white dwarf stage. A large UV field arises from the central (much hotter) star and at the same time, high velocity winds appear and interact with the AGB remnant (Cernicharo et al. 1989; Kwok 2000). Objects in this phase are called proto-planetary nebulae (hereafter {\\em PPNe}). During this period (1000$-$2000 years; Bujarrabal et al. 2001), the almost spherical symmetry typical of  the AGB stage disappears and is replaced by a complicated geometrical structure (most PPNe are elliptical, bipolar or quadrupolar, Zuckerman \\& Aller 1986; Frank \\etal 1993). The chemical composition is also strongly affected by these processes: O-bearing molecules can be formed in C-rich objects (Herpin \\& Cernicharo 2000), and complex organic molecules, which are not observed in IRC+10216, are efficiently produced (Cernicharo \\etal 2001 a\\&b).  In order to better understand this evolution, we present in this paper a comparative study of the millimeter, submillimeter and far-IR CO line emission from 3 objects representing different stages of this fast transition: CRL 2688, a very young PPN, CRL 618, a PPN, and NGC 7027, a young planetary nebula (PN). The different excitation coSnditions of the observed CO lines allow to probe different layers in the CSEs. In particular, we study the CSE remnant wind and the higher velocity winds. The CSE component is the \"normal\" AGB wind, and the wind component refers to a higher velocity wind which is probably not present during the AGB phase and which contributes to changing the morphology from spherical to bi-polar or other non-spherical geometries. The observations are presented in section 2. The representativity of each selected source is discussed in section 3. Section 4 is devoted to the analysis of the spectra of IRC+10216 as a reference of the AGB phase. In section 5 we discuss the wind structure of the young PPN CRL 2688. Section 6 is devoted to CRL 618 and section 7 to the analysis of NGC 7027 observations. An overall comparative discussion of the four objects is then presented in section 8. ",
        "conclusions": "The study of the selected 3 objects in rapid evolution from the AGB to the PN stage clearly shows the crucial importance that (i) the interaction between fast and slow winds; and (ii) the increase of UV flux as the central object evolves toward the white dwarf state have on the chemistry. The spectroscopic evidence of the increasing UV flux comes from the appearance of atomic and ionic lines in CRL 618, which later dominate the far-IR spectrum of NGC 7027. Atomic and ionic lines tend to appear (at ISO's sensitivity) when the central star is hotter than 10000 K (Fong et al. 2001; Castro--Carrizo et al. 2001). On the other hand, the spectroscopic signature showing that shocks become less important as the evolution goes on, comes from the wind velocity decrease seen in the CO profiles (this must be confirmed by further studies on more objects and only refers to the neutral gas). The strongest shocks occur just after leaving the AGB when the central star is ejecting large amounts of material in a very fast wind (the case of CRL 2688). The AGB remnant envelope is being shocked by an inner, faster wind developed in the PPN stage (i.e., the 200 \\kms $\\;$ wind in CRL 618). Most of the \\COd, \\COt and HCN emission is produced in these shocks. % %The \\COt abundance remains quite stable according to the AGB phase. % The fast increase of the stellar temperature, will produce a new UV-dominated chemistry when T$_{eff}\\geq$30000 K. These new conditions (UV photons and shocks) will deeply modify the constitution of the inner parts of the envelope. Indeed, in CRL 618, O-bearing molecules (\\water and OH, Herpin and Cernicharo 2000) and relatively complex organic molecules (Cernicharo et al., 2001a\\&b) appear. Furthermore, most of the  CO and HCN will be entirely reprocessed in the PDR leading to strong HNC emission. At this point, CO and [O$\\small I$] atomic lines are the dominant coolants.  As the star reaches the PN stage, the strong fast molecular winds have disappeared, and slow expanding layers constitute the PN envelope around a large and hot atomic region. Most of the {\\em old} AGB material has been reprocessed. The spectrum is now dominated by atomic and ionic lines. New species such as CH$^{+}$ and CH appear. There is only weak HCN emission, as the molecules may have been broken into H and CN. More interesting is the disappearance of \\water, which has probably also been reprocessed, and is only a relatively abundant molecule in the intermediate C-rich PPN stage. \\\\   {\\it Acknowledgments} We thank Spanish DGES, CICYT, and PNIE for funding support for this research under grants PB96-0883, ESP98-1351E and PANAYA2000-1784. JRG acknowledges \\textit{UAM} for a pre-doctoral fellowship. JRP also acknowledges further support for his research from NSF grants AST99-80846 (CSO operations) and ATM96-16766. We thank Dr. V. Bujarrabal for useful comments and suggestions. We thank the anonymous referee for his useful comments which result in an important improvement of this paper."
    },
    "0206/astro-ph0206477_arXiv.txt": {
        "abstract": "\\noindent  We present combined H$\\alpha$+HI rotation curves for a sample of spiral galaxies. Most of the velocity profiles (spectra at single points) in these galaxies are asymmetric, preventing the use of standard methods like the first moment analysis and the single Gaussian fitting. We thus propose a method similar to the Envelope-Tracing method (Sofue \\& Rubin 2001) to analyse those profiles from which we obtain HI rotation curves in good agreement with the H$\\alpha$ rotation curves. These final rotation curves provide the required high resolution in the inner parts of the galaxies, but also extend out to typically 2-3 R$_{\\rm opt}$. They will hence allow us to investigate the distribution of dark matter. ",
        "introduction": "The study of rotation curves of spiral galaxies provides the best evidence for dark matter on galactic scales. However, important properties like the shape of the dark halo, the distribution of dark matter and the relative importance of dark and luminous matter are still questions under debate. Recent results from the analysis of rotation curves (Borriello \\& Salucci 2001 for normal spirals, de Blok et al. 2002 for low surface brightness galaxies) seem to favour dark halos with constant density cores. However the debate is still going on as to whether the available data have sufficient quality to constrain the distribution of dark matter (Primack 2002). For this reason high-quality data with high resolution and large spatial extension are needed. This can be ideally achieved by combining H$\\alpha$ rotation curves (for the high resolution) and HI rotation curves (for the extension to large galactocentric radii). In particular, the high resolution in the inner parts is crucial to deduce the shape of the density profile (Blais-Ouellette et al. 2001), while the extension to large radii is needed to constrain the size of the possible core (Borriello \\& Salucci 2001). All these considerations can be done for the local Universe; the Atacama Large Millimeter Array (ALMA) will enable us to investigate the structure and evolution of dark matter halos at high redshifts. ",
        "conclusions": ""
    },
    "0206/astro-ph0206127_arXiv.txt": {
        "abstract": "The X-ray populations of Local group galaxies have been classified in detail by Einstein, ROSAT and ASCA revealing a mix of binaries, supernova remnants and HII regions.  However, these observatories were unable to resolve X-ray sources in galaxies beyond the local group.  With \\chandra's exquisite spatial resolution we are able to resolve sources in a sample of nearby galaxies.  We show that there are highly significant differences in the X-ray colors  of sources in bulge and disk systems.  In particular, we find that there is a population of X-ray soft, faint sources in disk galaxies not seen in bulges, and a smaller population of hard sources also seen preferentially in disk systems.  These differences can  be used as a basis to classify sources as low and high mass X-ray binaries, supernova remnants and supersoft sources. We suggest that the soft sources seen preferentially in disks are probably dominated by supernova remnants, although we cannot rule out the possibility that they are a new population of absorbed, faint, supersoft accretion sources associated with the young stellar population.  The hard sources seen in disks but not bulges we identify as high mass X-ray binaries.   While it is impossible to classify any individual source on the basis of X-ray color alone, the observations presented here suggest that  it is possible to separate sources into groups dominated by one or two source types.  This classification scheme is likely to be very useful in population studies, where it is crucial to distinguish between different classes of objects. ",
        "introduction": "\\label{intro} Our galaxy  and other nearby galaxies contain a multitude of X-ray sources \\citep{fab89}.  Observations of sources in the Milky Way  and  Local Group galaxies show that many of the brightest X-ray sources are  high and low mass  X-ray binaries (HMXB, LMXB), but there are also sources associated with supernova remnants (SNR) and HII regions. For example, the ROSAT survey of M31 \\citep{supper01} revealed a total of 560 X-ray sources, 16 of which have been optically identified with supernova remnants and 33 with globular clusters.  Most of the bulge sources in M31 have been identified on the basis of their X-ray spectra as low mass X-ray binaries  \\citep{trinch99,primini93}.  Studies of X-ray populations in a wide range of galaxies is potentially very valuable.  For example, since X-ray binaries are products of the end point of stellar evolution, the X-ray binary population  should be related to the history of star formation in each galaxy \\citep{wu01,rek02_1}.  However, the resolution of {\\it Einstein}, ROSAT and ASCA was not sufficient to allow identification of sources in galaxies beyond the Local Group. \\chandra\\ has the angular resolution necessary to study the X-ray source population in more distant galaxies for the first time \\citep{blanton01,tenn01,irwin02,fab01,soria02}.    In this paper we compare X-ray populations of a sample of disk and bulge galaxies.  We show that there is a highly significant difference in the X-ray colors of sources in bulge and disk systems.  Disk galaxies have a population of soft sources and a population of hard sources not seen in bulges.  X-ray colors are known to be a sensitive discriminator of source type \\citep{white84,haberl00,yok00,sasakil00} and we interpret these differences as differences in the source population.  Furthermore, we suggest that X-ray colors derived from \\chandra\\ observations can be used to identify different classes of X-ray sources in nearby galaxies.  In Section~\\ref{anal} we describe our data analysis. In Section~\\ref{sec:b_and_d} we show the difference in the distribution of X-ray colors in bulge and disk systems and in Section~\\ref{sec:id}  we suggest a classification scheme based on colors. In Section~\\ref{sec:abs} we discuss the effects of absorption on source classification, and in Section~\\ref{sec:LVSD} we discuss the spatial distribution, luminosities and variability characteristics of the different classes of sources.  Finally in Section~\\ref{sec:s_and_c} we summarize our results and give suggestions for further work. ",
        "conclusions": "\\label{sec:s_and_c}  In this paper, we show that there is a highly significant difference in the X-ray colors of sources in bulge and disk systems.  Disk galaxies have an additional population of soft X-ray sources and a scattering of hard sources not seen in bulge systems.  The hard disk sources are probably HMXBs.  The soft disk sources are probably dominated by SNR, but we cannot rule out a contribution from a new population of soft, faint, absorbed accretion sources associated with the young stellar population.  The differences in X-ray colors of sources in nearby disk and bulge galaxies can be used as a starting point for source classification, with the color-color diagram approximately separating SNR, LMXB and HMXB sources.  These conclusions are strengthened by the location of known thermal SNRs and binary pulsars in the X-ray color-color diagram.  The luminosities of sources identified as SNR and HMXBs are consistent with what is observed in the Galaxy \\citep{grimm01}.  The variability characteristics of different classes of sources should provide additional constraints on their nature. Accreting binaries are likely to vary stochastically while the flux from thermal SNR should remain constant.  Unfortunately, although there are two \\chandra\\ observations for both disk galaxies (M101 and M83) the second observations are not deep enough to effectively constrain the variability of the soft sources identified as SNR.  It is not possible to identify a source on the basis of X-ray color alone.  However, separating sources on the basis of X-ray colors will be very valuable for population studies. Although X-ray colors provide an excellent starting point for the classification of sources in nearby galaxies, there is some overlap in source type in the X-ray color-color diagram.  For example, supernova remnants which are very highly absorbed will move vertically upwards into the ``LMXB'' part of the color-color diagram.  The population of soft sources may contain some supersoft sources or X-ray binaries with soft spectra.  Therefore other information such as source variability, luminosity, and optical counterparts as well as position within the galaxy  must be used to complete the classification process. Deeper \\chandra\\ images of nearby galaxies are required to investigate the variability and spectral  properties of the different source classes outlined here."
    },
    "0206/astro-ph0206311_arXiv.txt": {
        "abstract": "Using recent dust continuum data, we generate the intrinsic ellipticity distribution of dense, starless molecular cloud cores. Under the hypothesis that the cores are all either oblate or prolate randomly-oriented spheroids, we show that a satisfactory fit to observations can be obtained with a gaussian prolate distribution having a mean intrinsic axis ratio of 0.54. Further, we show that correlations exist between the apparent axis ratio and both the peak intensity and total flux density of emission from the cores, the sign of which again favours the prolate hypothesis. The latter result shows that the mass of a given core depends on its intrinsic ellipticity. Monte Carlo simulations are performed to find the best-fit power law of this dependence. Finally, we show how these results are consistent with an evolutionary scenario leading from filamentary parent clouds to increasingly massive, condensed, and roughly spherical embedded cores. ",
        "introduction": "\\label{sec-intro} Little is known with certainty about the intrinsic shapes of the dense molecular clouds that give birth to stars. The present situation is reminiscent of the analogous study of elliptical galaxies some thirty years ago, when relatively few tests had been performed on limited datasets. In that field, attempts to model the observed distributions of ellipticity, surface brightness, and velocity dispersion with axisymmetric objects (e.g., oblate or prolate spheroids) met with limited success (Merritt 1982). Additional kinematical data and $N$--body simulations have led to the conclusion that many ellipticals are, in fact, triaxial (Merrifield and Binney 1998).  While from a modelling perspective our understanding of molecular cloud shapes lags behind that of stellar systems, much can be learned by applying the same methods in this new arena. At the same time, it is important to keep in mind certain salient differences between the two contexts, aside from the obvious difference of scale. First and foremost, gaseous self-gravitating clouds are not collisionless systems. Agents such as thermal instability, pressure gradients, and magnetic fields, more or less unique to the ISM context, have all been shown to be important in various physical regimes. Second, studies of the most centrally condensed cores of molecular clouds reveal them to be not the isolated (or infinite) balls of dense gas considered by simple theory, but rather the lowest rung in a hierarchy of structure beginning with parsec-scale entities (giant molecular clouds). This aspect needs to be considered when theory is brought to bear upon intrinsic core properties.  Clues to core structure are beginning to be extracted from morphological studies. Early analyses were hampered by datasets of limited size, which reduced the statistical significance of the conclusions (David \\& Verschueren 1987; Myers et al.\\ 1991; Ryden 1996). Nevertheless, under the hypothesis that each object is a spheroid randomly-oriented to the line of sight, each of these studies concluded that cores were more likely to be intrinsically prolate than oblate. The recent study of Jones, Basu, \\& Dubinski (2001) analyzed the largest dataset to date: 264 ammonia cores compiled by Jijina, Myers, \\& Adams (1999). By showing that the best-fit probability distributions of prolate and oblate spheroids became negative near $p=1$, these authors rejected the hypothesis of axisymmetry altogether. A closer examination of the analysis technique, however, gives a likely explanation of their results (see below \\S \\ref{sec-Nvsp}).  All of the aforementioned analyses utilized but one observational diagnostic of core shape: the projected axis ratio, or ellipticity. However, as remarked by Fleck (1992), the {\\it mean} value of this quantity cannot be used to distinguish one spheroidal shape over another. Moreover, as pointed out by Binney \\& de Vaucouleurs (1981) and Ryden (1996), the measurement of axis ratios is subject to various systematic biases that can affect the overall distribution and subsequent analysis. In fact, other physical properties---such as column density, velocity dispersion, and mass---have influenced theoretical models of cloud cores far more than have the perceived ellipticities. We feel that to abandon the hypothesis of axisymmetry in lieu of examining these diagnostics is premature, especially given the lack of a physical basis for triaxiality in the ISM context (akin to the anisotropic velocity distribution of stars in elliptical galaxies).  In \\S \\ref{sec-observ} and \\ref{sec-intrins}, we employ a number of tests to discriminate between the oblate and prolate spheroidal hypotheses. Two of the tests are new in the ISM context, and the other---the distribution of projected axis ratios---has not previously been applied to the dust continuum data on which we base our analysis (however, subsets of these data were recently analyzed by Jones \\& Basu (2002), under the oblate and triaxial hypotheses only). Further, by simulating the observed sample using distributions of model cores of both types, we are able to place rather stringent constraints on one intrinsic property, the variation of polar intensity with intrinsic ellipticity. All of the tests are independent of distance, still a very uncertain quantity for these objects (see, e.g., Launhardt \\& Henning 1997).  Finally, in \\S \\ref{sec-discuss} we interpret our results in the framework of an evolutionary sequence of core shapes as a function of time. In this picture, core morphology at early times is primarily determined by the nature of the surrounding nonisotropic mass distribution. As the core grows in mass, it approaches a spherical shape consistent with the dominance of self-gravity. As we shall see, this behavior is consistent with the present observational picture. ",
        "conclusions": "\\label{sec-conc} Previous studies of dense core morphology, nearly all of which used molecular line data, gave varied results as to the intrinsic shapes of the cores. We have employed recent dust continuum datasets and additional methods of analysis in an effort to clarify the situation. The main conclusions of this work are as follows: \\vskip 0.1 cm \\noi (1) The observed distribution of core ellipticities in the combined continuum sample is well fit by a gaussian distribution of intrinsically prolate objects with mean ellipticity $\\lb q \\rb \\approx 0.5 \\pm 0.2$. \\vskip 0.1 cm \\noi (2) In the M01 sample, the peak intensity $I$ is positively correlated with the apparent ellipticity $p$, with a slope in log $I$ vs.\\ log $p$ of $0.50 \\pm 0.18$. This slope is shallower than the $+1$ value expected for an ensemble of randomly-oriented prolate spheroids, each having constant polar intensity $I_p$. \\vskip 0.1 cm \\noi (3) In the same sample, an equally significant correlation (slope $= 0.65 \\pm 0.21$) is observed between the log of the total flux density $S$ and log $p$. This shows that both $I_p$ and the mass of a given core depend on its intrinsic ellipticity, $q$. \\vskip 0.1 cm \\noi (4) Under the assumption that $I_p = q^m I_0,~I_0 =$ constant, Monte Carlo simulations were used to find the value of $m$---and the spheroidal shape---that best fits the observed correlation in (log $I$, log $p$). The observed slope and rank C.C.\\ are best fit by a prolate ensemble with $m = -0.60$ and an intrinsic scatter in log $I_0$ = 0.15. No satisfactory fits for oblate spheroids were found. \\vskip 0.1 cm \\noi (5) The relation $I_p = q^{-0.6} S_0$ was shown to agree (for $q \\gae 0.25$) with the expected polar column density in the embedded prolate equilibrium sequence of Curry (2000). \\\\  The chief limitation of the present work is the relatively small number of objects in the M01 sample, and the consequent weakening of the statistical results so obtained. Also, the fact that all of the cores come from mainly filamentary structures found in three regions may be considered a bias. However, if these cores are truly pre-stellar (as suggested by their number distribution as a function of mass; see \\S \\ref{sec-intro}), then this mode of condensation may in fact be reasonably representative of low-- and intermediate--mass star formation. The possibility that at least some of these highly embedded objects are aligned with the major axes of their parent filaments is real, and may alter the results presented here in certain respects (\\S \\ref{sec-orient}). The addition of new dust continuum data will allow more definitive conclusions to be drawn on each of these key points, and will no doubt aid in the formulation of more sophisticated theoretical models of embedded cloud equilibria. \\\\\\\\ \\noi I am grateful to F.\\ Motte for providing her previously unpublished peak fluxes, and thank Carol Jones for preliminary discussions about this work. Chris McKee, Steve Stahler, and an anonymous referee offered valuable comments that helped improve the original manuscript."
    },
    "0206/astro-ph0206061_arXiv.txt": {
        "abstract": "This paper discusses a quasi-static evolution of a force-free magnetic field under slow sheared footpoint motions on the plasma's boundary, an important problem with applications to the solar and accretion disk coronae. The main qualitative features of the evolution (such as field-line expansion and opening) are considered and a comparison is made between two different geometrical settings: the Cartesian case with translational symmetry along a straight line, and the axisymmetric case with axial symmetry around the rotation axis. The main question addressed in the paper is whether a continuous sequence of force-free equilibria describes the evolution at arbitrarily large values of the footpoint displacement or the sequence ends abruptly and the system exhibits a loss of equilibrium at a finite footpoint displacement. After a formal description of the problem, a review/discussion of the extensive previous work on the subject is given. After that, a series of simple scaling-type arguments, explaining the key essential reason for the main qualitative difference between the two geometry types, is presented. It is found that, in the Cartesian case, force-free equilibria exist at arbitrarily large values of shear and the field approaches the open state only at infinite shear, whereas in the axisymmetric case the field opens up already at a finite shear. ",
        "introduction": "\\label{sec-intro}  In studies of force-free coronal magnetic fields in solar physics, as well as in a closely related and essentially very similar problem of accretion disk magnetospheres, there has been some controversy regarding the issue of the {\\it loss of equilibrium}. This controversy has arisen from the problem of finding a continuous sequence of force-free equilibria in the corona, invoked to represent a time evolution of the coronal magnetic field under slow plasma motions in the Sun's photosphere. Indeed, the footpoints of the magnetic field lines are frozen into the photosphere and hence the photospheric motions lead to continuous shearing of the magnetic field. Under the assumption of ideal magnetohydrodynamics (MHD), if these motions are much slower than the Alfv{\\'e}n velocity in the corona (an assumption justified by a very low plasma density in the corona), the coronal magnetic field progresses through a sequence of equilibria. An important aspect of the problem is that the domain under consideration is infinite, so that the field lines can expand freely into space, instead of being confined to a finite-size box.  When trying to build a theoretical model of this process, one typically starts with a potential (no shear) field, and then gradually increases the shear. This initial potential field is taken to be closed, which means that both footpoints of each field line lie on the surface, i.e., no field line extends to infinity. The critical question then is whether one should be able to find a force-free equilibrium configuration (with the same topology as that of the original potential field) as the shear is increased indefinitely, or one should reach a certain critical point, beyond which no same-topology equilibrium solutions can be found (loss of equilibrium).  Even though this question has first been tackled in the context of the solar corona, a very similar process has also been investigated in the context of accretion disk coronae, where the sheared footpoint motion arises naturally from the differential rotation of a Keplerian disk or from the relative star--disk rotation (see van~Ballegooijen 1994; Goodson et al. 1999; Uzdensky et al. 2002; Lovelace et al. 1995; Uzdensky 2002).  In either context, some simplifying assumptions are usually made in order to make the problem tractable. One of the most important is the assumption that there is one ignorable direction on the photospheric surface, i.e., a direction along which the fields are constant. There are two different symmetry classes that are most often studied:  1) {\\it Cartesian (or plane) geometry} (see Fig.~\\ref{fig-cartesian}), the photosphere being an infinite plane and the corona --- a half-space above this plane. The field line topology is that of a straight line dipole placed on or somewhere beneath the plane. The footpoints are displaced along the line dipole axis (also called the polarity inversion line) and the system possess translational symmetry along this axis. This problem is usually studied in Cartesian coordinates, with the ignorable direction (the line dipole axis) denoted by, say, $z$ and the direction perpendicular to the plane by~$y$. There have been extensive analytical (Low 1977, 1982, 1990; Birn et al. 1978; Priest \\& Milne 1980; Birn \\& Schindler 1981; Aly 1984, 1985, 1990, 1993, 1994; Priest \\& Forbes 1990) and numerical studies of this problem. The latter can be subdivided further into numerical computations of sequences of force-free equilibria (Sturrock \\& Woodbury 1967; Jockers 1978; Klimchuk et al. 1988; Klimchuk \\& Sturrock 1989; Finn \\& Chen 1990; Wolfson \\& Verma 1991) and full MHD numerical simulations (e.g., Biskamp \\& Welter 1989; Amari et al. 1996a).  2) {\\it axisymmetric geometry}, with axial (or cylindrical) symmetry around the $z$ axis ($\\phi$-direction). This case is usually treated in either cylindrical ($\\rho,z,\\phi$) or spherical ($r,\\theta,\\phi$) coordinates. The footpoints rotate in the azimuthal (or toroidal) direction~$\\phi$. The problem has been considered both analytically (Aly 1984, 1991, 1993, 1995; Low 1986; van~Ballegooijen 1994; Lynden-Bell \\& Boily 1994; Sturrock et al. 1995; Wolfson 1995; Uzdensky et al. 2002; Uzdensky 2002) and numerically (Barnes and Sturrock 1972; Yang et al. 1986; Porter et al. 1992; Wolfson \\& Low 1992; Roumeliotis et al. 1994; Miki{\\'c} and Linker 1994; Uzdensky et al. 2002).  One should be aware that the are actually two distinct geometrical settings in the axisymmetric case. One (which we shall call {\\it spherical geometry}) is where the domain if interest is the outside of a differentially rotating sphere, with all the footpoints fixed on the surface of the sphere (studied, for example, by Low 1986; Wolfson \\& Low 1992; Miki{\\'c} \\& Linker 1994; Roumeliotis et al. 1994; Wolfson 1995; Aly 1995; Sturrock et al. 1995). This case is usually considered in the context of solar corona. The field topology here is that of a point dipole placed inside the sphere (see Fig.~\\ref{fig-spherical}). The other case, superficially similar to the Cartesian one, is {\\it cylindrical geometry}, where the domain of interest is the half-space above an infinite plane on which all the footpoints are fixed (considered by Barnes and Sturrock 1972; Yang et al. 1986; Porter et al. 1992; Lynden-Bell \\& Boily 1994; Sturrock et al. 1995 among others). The field topology in the cylindrical case can be visualized by, for example, placing a ring dipole on a plane surface or by putting a point dipole (with its axis being perpendicular to the plane) underneath the plane surface (see Fig.~\\ref{fig-cylindrical}). This geometry is relevant to both the solar and accretion disk studies. Finally, a magnetically-linked star--disk system involves a combination of these two settings (van~Ballegooijen 1994; Goodson et al. 1999; Uzdensky et al. 2002). ",
        "conclusions": "\\label{sec-summary}  In this paper we investigated the question of finite-time opening of a force-free magnetic field evolving quasi-statically under slow footpoint motions on the boundary. This problem is of great importance in studies of both the solar corona and accretion disk magnetospheres. We started (in \\S~\\ref{sec-intro}) by discussing two principal classes of the problem's geometry that are most frequently considered: the Cartesian geometry with the translational symmetry along a straight line, and the axisymmetric geometry with the symmetry with respect to rotations around an axis. This latter class includes both the spherical geometry, where the domain of interest is the outside of a sphere, and the cylindrical geometry, where the domain of interest is the half-space above a plane. In \\S~\\ref{sec-2approaches} we introduced the set of equations governing the force-free evolution in these two geometrical settings. After that we gave a review of the existing literature focusing on various approaches and ways to describe the footpoint shearing responsible for driving the evolution. This followed by our compilation of the most commonly accepted (in our opinion) scenario for both the Cartesian and the axisymmetric cases. The main aspects of the evolution can be described as follows (see \\S~\\ref{sec-2approaches}). In both cases the evolution consists of two phases. During the first phase the shape of the flux surfaces changes very little, while the magnetic field's component in the ignorable direction (which in the simplistic setting considered here coincides with the direction of the footpoint motion) grows monotonically in time. After some finite shearing, however, this ``toroidal'' component stops growing; the system enters the second phase, during which the ``poloidal'' (i.e., perpendicular to the ignorable direction) field starts expanding rapidly, while the toroidal field on the photospheric surface decreases. Eventually the system approaches the open (or partially-open) field configuration. The main difference between the two types of geometry is that in the Cartesian case the opening is achieved only asymptotically in the limit of infinite shear, whereas in the axisymmetric case it is most likely achieved at some finite shear (finite-time field opening).  Finally, in \\S~\\ref{sec-finite-time-opening} we presented a series of several simple physical arguments invoked to explain the geometrical origin of this most important qualitative difference in the character of the field-line expansion and opening process between the two cases. The basic idea of these arguments can be described as follows. The field-line inflation is driven by the toroidal field's pressure, which in a force-free equilibrium is balanced by the tension of the poloidal field. Hence, in the outer parts of a strongly-expanded field line the toroidal field's strength should in some sense be comparable with that of the poloidal field. This, in turn requires that the toroidal extent of the expanded field lines be comparable with their very large poloidal extent. In the Cartesian case, this condition automatically leads to the necessity of having a comparably large footpoint displacement. Therefore, the opening of the field cannot occur at any finite shear in this case. In the axisymmetric case, on the other hand, a very large toroidal extent of a field line's outermost portion can be reached at a finite footpoint rotation angle, which leads to a finite-twist field opening.  Now let us discuss some limitations of the simple picture described in the present  paper. One of the implied assumptions in the paper is that the opening is approached via a continuous sequence of stable equilibria. In a very interesting alternative scenario presented (in spherical geometry) by Wolfson \\& Low (1992), a sudden transition to a {\\it partially-open state} is suggested to take place as soon as the energy of the twisted but still closed magnetic field configuration exceeds that of the partially open field (see also Low 1990). Note that according to the Aly--Sturrock conjecture (Aly 1984, 1991; Sturrock 1991) the energy of the closed field can never exceed that of the {\\it totally open} field. Therefore, only partial opening can be achieved by such a sudden eruption. Also note that this scenario can only work in the axisymmetric geometry because in the Cartesian case the energy (per unit length in the ignorable direction) of even partially open field is infinite and hence can never be exceeded by that of any closed-field configuration.  In addition, it is worth mentioning that the discussion in this paper assumed that ideal MHD is valid throughout the entire evolution. If finite resistivity exists in the system, then reconnection may take place at a finite shear, leading to a change in the field topology, e.g., to the formation of a plasmoid (e.g., Amari~et~al. 1996a). This reconnection process is likely to take place soon after the energy of the sheared arcade exceeds the energy of the complex-topology configuration with the plasmoid (Aly 1990, 1991, 1993). This phase of the evolution may be very rapid and dynamic, perhaps characterized by the plasmoid ejection (Aly 1993; Amari~et~al. 1996a), hence providing a possible mechanism for coronal mass ejections.  Finally, it is important to acknowledge that this paper deals exclusively with the question of {\\it existence} of equilibria; the very important issue of {\\it stability} of these equilibria is completely left out. The main reason for this is that one cannot analyze the stability of solutions before establishing their existence and, mathematically, existence studies can be done completely independently of stability studies. From a practical perspective, however, it is clear that a quasi-static evolution is physically meaningful only if the sequence is made of continuous stable equilibria. It is possible, for example, that the first branch of the evolution is always stable, while the solutions on the second branch are unstable' the transition state between the two branches is then a marginally stable state. This is in fact the essence of Low's (1990) suggestion that the gradual quasi-static sequence of equilibria would end at this marginally stable bifurcation state. In particular, he argued that the presence of even a small plasma pressure may render this state unstable, leading to a subsequent violent phase of evolution. Thus, a stability analysis of strongly-inflated force-free equilibria is of crucial importance and definitely presents the next logical step in the study of sheared force-free systems, but it falls outside of the scope of our present work.  I am grateful to J.~J.~Aly and B.C.~Low for their interesting and useful comments. I would like to acknowledge the support by the NSF grant NSF-PHY99-07949."
    },
    "0206/gr-qc0206026_arXiv.txt": {
        "abstract": "Assuming a flat Friedmann-Robertson-Walker cosmology with a single perfect fluid, we propose a pressure-density ratio that evolves as a specific universal function of the scale parameter. We show that such a ratio can indeed be consistent with several observational constraints including those pertaining to late-time accelerated expansion. Generic dynamical scalar field models of Dark energy (with quadratic kinetic terms in their Lagrange density) are shown to be in accord with the proposed equation-of-state ratio, provided the current matter density parameter $\\Omega_{m0} < 0.23$ - a value {\\it not} in agreement with recent measurements. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0206/astro-ph0206257_arXiv.txt": {
        "abstract": "We have undertaken quantitative analysis of four LMC and SMC O4--9.7 extreme supergiants using far-ultraviolet {\\fuse}, ultraviolet {\\iue}/{\\hst} and optical {\\vlt} UVES spectroscopy. Extended, non-LTE model atmospheres that allow for the consistent treatment of line blanketing (Hillier \\& Miller 1998) are used to analyse wind and photospheric spectral features simultaneously. Using H$\\alpha$ to constrain $\\dot{M}$, He\\,{\\sc i-ii} photospheric lines reveal stellar temperatures which are systematically (5--7.5kK) and substantially (15--20\\%) lower than previously derived from unblanketed, plane-parallel, non-LTE photospheric  studies. We have confidence in these revisions, since derived temperatures generally yield consistent fits across the entire $\\lambda\\lambda$912--7000\\AA\\ observed spectral range. In particular, we are able to resolve the UV-optical temperature discrepancy identified for \\oseven\\ (O7\\,Iaf$^+$) in the SMC by Fullerton et al. (2000).  The temperature and abundance sensitivity of far-UV, UV and optical lines is discussed. `Of' classification criteria are directly linked to (strong) nitrogen enrichment (via N\\,{\\sc iii} $\\lambda$4097) and (weak) carbon depletion (via C\\,{\\sc iii} $\\lambda\\lambda$4647-51), providing evidence for mixing of unprocessed and CNO processed material at their stellar surfaces. Oxygen abundances are more difficult to constrain, except via O\\,{\\sc ii} lines in the O9.7 supergiant for which it is also found to be somewhat depleted. Unfortunately, He/H is very difficult to determine in individual O supergiants, due to uncertainties in microturbulence and the atmospheric scale height. The effect of wind clumping is also investigated, for which P\\,{\\sc v} $\\lambda\\lambda$1118--28 potentially provides a useful diagnostic in O-star winds, unless phosphorus can be independently demonstrated to be underabundant relative to other heavy elements. Revised stellar properties affect existing calibrations of (i) Lyman continuum photons -- a factor of two lower for the O4 supergiant; and (ii) kinetic energy released into the ISM by O supergiants. Our results also have importance for the calibration of the wind momentum-luminosity relationship for OB stars, particularly since the stars studied here are amongst the visually brightest OB stars in external galaxies. ",
        "introduction": "The existence of winds in O-type stars has been established since the 1960's, when the first rocket-ultraviolet (UV) observations revealed the characteristic resonance line P~Cygni signatures of mass loss (Morton 1967). Far-ultraviolet (FUV) spectroscopy of O-type stars with {\\it Copernicus}, the {\\it Hopkins Ultraviolet Telescope} and {\\it ORFEUS} missions revealed many additional stellar-wind features. The launch of the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\fuse}) telescope (Moos et al. 2000) has provided a new opportunity to study a wide variety of OB stars spanning a range of metallicities, $Z$, at high spectral-resolution.  Of primary interest is the dependence of mass-loss rates for luminous OB stars on $Z$. Theoretically, the strength of radiatively driven O-star winds is predicted to depend on metallicity, $Z$, as $\\dot{M} \\propto Z^{0.5-0.7}$ (Kudritzki, Pauldrach \\& Puls 1987; Vink, de Koter \\& Lamers 2001). Observationally, the principal method of deriving mass-loss properties of O stars has been via radio observations for stars within a few kpc, or H$\\alpha$ observations more generally (Puls et al. 1996). {\\fuse} revives the possibility of using UV  resonance lines to determine mass-loss rates empirically.  Observational results from H$\\alpha$ have been combined with theoretical predictions to generate the so-called Wind-Momentum-Luminosity Relationship (WLR; Kudritzki \\& Puls 2000), which can be used to determine extragalactic distances. To date, such studies almost exclusively employed plane-parallel, non-LTE models for determinations of temperature (using photospheric He lines), and separately use wind models to determine mass-loss rates (using H$\\alpha$ or UV wind profiles). Recent examples of the latter approach include Pauldrach et al. (1994) and Pauldrach, Hoffman \\& Lennon (2001). Clearly, the underlying assumption in such studies is that the effects of winds on photospheric optical lines is negligible -- see, however, Gabler et al. (1989) and Schaerer \\& Schmutz (1994). Generally either line blanketed or spherically extended model atmospheres have been employed (e.g. Herrero, Puls \\& Villamariz 2000). Rarely have both effects been considered simultaneously.  Since OB supergiants have an enormous impact on the chemical and dynamical evolution of their environments, their properties are of considerable interest. Recently, Fullerton et al. (2000; hereafter Paper~I) presented a study of {\\oseven} (O7\\,Iaf$^{+}$, SMC) and Sk-67$^{\\circ}$ 111 (O6\\,Ia(n)fp var\\footnote{Classification revised recently by Walborn et al. 2002a}, LMC) based on {\\fuse} spectroscopy. Non-LTE wind models allowing for line blanketing and an expanding atmosphere (Pauldrach et al. 2001) were used to constrain their stellar temperatures to $T_{\\rm eff}\\sim$32kK based on {\\fuse} FUV wind profiles. In contrast, recent non-LTE, plane-parallel, hydrostatic  studies of {\\oseven} (e.g. Puls et al. 1996), derived a substantially higher stellar temperature of $T_{\\rm eff}\\sim$38kK from optical photospheric lines. This represents an important discrepancy with respect to derived bolometric luminosities ($\\propto T_{\\rm eff}^{4}$), and so indirectly affects the calibration of the WLR.  In this paper, we investigate the properties of a small sample of extreme OB supergiants in the Magellanic Clouds to verify whether the supergiants studied in Paper~I are typical, and if so, how these discrepancies may be resolved. These stars were selected to cover a wide range of spectral types, with the additional criteria that (i) they should have low interstellar H$_2$ column densities to minimize contamination in the {\\fuse} region; and (ii) they should have small projected rotational velocities. We supplement the {\\fuse} spectroscopy with {\\iue} and {\\hst} UV spectroscopy, together with ground-based optical observations in order to help resolve previously conflicting determinations of stellar temperature and mass-loss rate. To date, UV studies of OB stars generally adopt temperatures from optical plane-parallel analyses (e.g. Haser et al. 1998). The sole exception was HD\\,93129A (O2~If*, Walborn et al. 2002b) for which Taresch et al. (1997) consistently analysed its optical, UV, and FUV spectrum.  The paper is structured as follows. Observations of the four program O supergiants are presented in \\S~\\ref{obs}, followed by a determination of stellar parameters using standard plane-parallel methods in \\S~\\ref{hydro}. Spherical, line-blanketed models are utilised in \\S~\\ref{sect4}, revealing a substantial revision in stellar properties. Abundances, including CNO elements are discussed in \\S~\\ref{abundances}, whilst clumping in O supergiants is considered in \\S~\\ref{clumping}. Finally, our conclusions are reached in \\S~\\ref{conclusions}. ",
        "conclusions": "Our analysis of optical H-He wind and photospheric profiles of luminous O supergiants, using line-blanketed, extended model atmospheres, reveal systematically ($\\sim$15--20\\%) lower temperatures than plane-parallel results based solely on optical H-He photospheric lines (e.g. Herrero et al. 1992). Our results are supported by UV {\\hst} and especially FUV {\\fuse} spectroscopy of metal wind lines and photospheric iron lines in these stars. Initial {\\fuse} datasets of {\\oseven} raised questions about the validity of temperatures derived from plane-parallel O supergiant models (Paper~I), which unified model atmospheres such as Hillier \\& Miller (1998) can now resolve.  Fig.~\\ref{osuper} compares temperatures for O supergiants derived here with those from previous compilations by B\\\"{o}hm-Vitense (1981), Schmidt-Kaler (1982), Howarth \\& Prinja (1989) and most recently by Vacca et al. (1996). Of these, one might expect Vacca et al. to be the closest match to the present results, yet the reverse is true, especially at the earliest subtypes. This can be understood from the fact that Vacca et al. included the most recent spectroscopic results, {\\it excluding} wind or line blanketing. Consequently, $\\zeta$ Pup (O4\\,I(n)f) was included in their calibration with $T_{\\rm eff}$(no wind blanketing)=46.5kK, rather than $T_{\\rm eff}$(wind blanketing)=42kK (Bohannan et al. 1986), although the `core-halo' approach was followed in their study. Test calculations carried out  for $\\zeta$ Pup, which consistently include sphericity {\\it and} line blanketing, indicate $T_{\\rm eff}\\sim39$kK, illustrating the additional effect that line blanketing plays.  More recently, Herrero et al. (2000) compared results of early O supergiants obtained with plane-parallel models that accounted for line blanketing with those obtained using spherical models for which blanketing was omitted. For HD\\,15570 (O4\\,If$^+$), Herrero et al. obtained $T_{\\rm eff}$(blanketed)=50kK versus $T_{\\rm eff}$(spherical)=42kK. Stellar temperature determinations of O supergiants undertaken during the past decade are summarised in Table~\\ref{t6}, sorted by the degree of complexity implemented; i.e., plane-parallel versus spherical, plus unblanketed versus line blanketed. From Table~\\ref{t6}, blanketing {\\it and} sphericity have rarely previously been simultaneously considered (Pauldrach et al. 1994, 2001), with only Taresch et al. (1997) successfully combining optical/UV/FUV diagnostics.  A critical revision to the temperature calibration of O supergiants clearly requires analysis of a substantially larger sample of targets, using the methods outlined here, which is currently underway. Revisions do not necessarily possess a strong metallicity dependence, since {\\oseven} (SMC) possesses a similar difference from standard calibrations to the three LMC stars studied here (Fig.~\\ref{osuper}). Wind strength is more critical -- one would expect the greatest deviations from conventional temperature scales for those stars with the highest wind densities, i.e. those with H$\\alpha$ and He\\,{\\sc ii} $\\lambda$4686 emission.  Similar calculations for O dwarfs -- also allowing for line blanketing -- obtained $\\sim$5\\% lower temperatures than unblanketed results (Herrero, Puls \\& Villamariz 2000; Martins et al. 2002). Independent observational evidence in favour of  (1--2kK) lower temperatures of O dwarfs may be drawn from studies of O-type  binaries (e.g. Harries, Hilditch \\& Hill 1998). The rarity of O supergiants within short-period eclipsing binary systems prevents direct determinations of radii, and thus temperatures, for O supergiants. The greater effect identified here for extreme O supergiants may reasonably be attributable to the effect of strong winds on the photospheric lines.  Our results argue for a substantial revision in stellar parameters for O supergiants versus those derived from previous spectroscopic techniques. For {\\ofour}, the reduction in temperature implies a decrease in luminosity from $\\sim$1.6$\\times 10^{6} L_{\\odot}$ to 1.0$\\times 10^{6} L_{\\odot}$. Such changes greatly affect the number of Lyman-ionizing photons emitted -- in the case of {\\ofour}, the standard $T_{\\rm eff}$-calibration would imply that its Lyman-ionizing output is a  factor of two higher than the 10$^{49.7}$ ph\\,s$^{-1}$ determined here. The ionizing flux of {\\ofour} below the He\\,{\\sc i} $\\lambda$504 edge is reduced by a factor of three to 10$^{49.0}$ ph\\,s$^{-1}$.  Similar changes are obtained for the other program stars. Fortunately, comparisons of nebular strengths in young, massive clusters with their constituent O stars are generally weighted towards main-sequence populations. Nevertheless, cases exist where individual O stars dominate H\\,{\\sc ii} regions (e.g. Oey et al. 2000) which would be dramatically affected by such large temperature changes.  Masses are more difficult to constrain, but the decrease in the spectroscopic luminosity and gravity also cause large differences from previous determinations. Consequently, evolutionary model comparisons with O supergiants carried out previously may not have been using the appropriate tracks in many cases. Although initial (and indeed current) rotational velocities for the program stars are not known, adopting $v_{\\rm init}$=300 km\\,s$^{-1}$ suggest an {\\it initial}  mass of 75$M_{\\odot}$ for {\\ofour} according to evolutionary models (Fig.~8 of Meynet \\&  Maeder 2000). In contrast, an initial mass in excess of 120$M_{\\odot}$ would be implied by models on the basis of its previously determined higher luminosity. Our results do not solve the long-standing `mass discrepancy' for O stars (Herrero et al. 1992) between evolutionary and spectroscopic mass determinations,  since the reduction in spectroscopic luminosity is also accompanied by a reduced mass. Again, using {\\ofour} as an example, the reduction in spectroscopic gravity from $\\log g\\sim$3.7 to $\\log g\\sim$3.35 implies a decrease from $\\sim70M_{\\odot}$ to $\\sim35M_{\\odot}$ in its {\\it current} mass.  The revision in luminosity also implies a different mass-loss rate. Table~\\ref{t4} shows a 35\\% reduction in $\\dot{M}$ for {\\ofour} (without accounting  for the  possibility that the wind is clumped). This revision affects the kinetic energy injected by individual stars or a young, massive cluster to the ISM, since this is more dependent on the O supergiants and WR properties than on the main sequence stars (e.g. Crowther \\& Dessart 1998).  These effects have major consequences for the Wind-Momentum-Luminosity Relationship (WLR). In Fig.~\\ref{momentum}, we present the WLR for Galactic (circles), LMC (squares) and SMC (triangles) luminous O stars from Puls et al. (1996, open symbols). Also shown is the form of the Wind-Momentum-Luminosity Relationship (dotted line) for Galactic O supergiants according to Kudritzki \\& Puls (2000). We have added current results for our sample of LMC and SMC targets (filled symbols), two of which are in common with Puls et al. (1996). Revisions in parameters are illustrated in these cases with arrows. Clearly, when the present results are extended to a larger sample of luminous OB stars, substantial revisions to the empirical relationship of Kudritzki \\& Puls (2000) are expected.   FUV {\\fuse} spectroscopy has proved to be  invaluable for the present analysis. In contrast with {\\hst}/{\\iue}, the availability of unsaturated resonance lines from dominant ionization stages of ``cosmically rare'' elements (e.g. P\\,{\\sc v} $\\lambda\\lambda$1118--28) represents a potentially exciting new diagnostic of wind clumping in  O stars. Elements such as S and P are especially useful since they do not change substantially dueing the evolution of a massive star, in sharp contrast to CNO. {\\fuse} provides many additional wind features, such that we no longer have access to only the saturated (model insensitive) C\\,{\\sc iv} $\\lambda\\lambda$1548--51 and X-ray influenced N\\,{\\sc v} $\\lambda\\lambda$1238--42 wind lines. Our ongoing {\\fuse} program will aim to undertake detailed modeling of UV wind lines for a large sample of O stars, covering a greater range  of spectral types and luminosity classes."
    },
    "0206/hep-th0206155.txt": {
        "abstract": "{ We propose a new approach for using the AdS/CFT correspondence to study quantum black hole physics. The black holes on a brane in an AdS$_{D+1}$ braneworld that solve the classical bulk equations are interpreted as duals of {\\it quantum-corrected} $D$-dimensional black holes, rather than classical ones, of a conformal field theory coupled to gravity. We check this explicitly in $D=3$ and $D=4$.  In $D=3$ we reinterpret the existing exact solutions on a flat membrane as states of the dual $2+1$ CFT. We show that states with a sufficiently large mass really are $2+1$ black holes where the quantum corrections dress the classical conical singularity with a horizon and censor it from the outside. On a negatively curved membrane, we reinterpret the classical bulk solutions as quantum-corrected BTZ black holes. In $D=4$ we argue that the bulk solution for the brane black hole should include a radiation component in order to describe a quantum-corrected black hole in the $3+1$ dual. Hawking radiation of the conformal field is then dual to classical gravitational bremsstrahlung in the AdS$_5$ bulk.}   % ",
        "introduction": "We propose here a connection between two seemingly unrelated problems in black hole theory: {\\it i)} the well-known problem of the backreaction from quantum effects on a black hole geometry, and {\\it ii)} the description of a black hole in an AdS braneworld, as in the Randall-Sundrum model with an infinite extra dimension, RS2 \\cite{rs}. Quantum fields in a black hole background lead to particle production and black hole evaporation via Hawking radiation \\cite{swh}. To leading order in perturbation theory, this yields an expectation value of the renormalized stress-energy tensor of quantum fields $\\left<T_{\\mu\\nu} \\right>$, which includes quantum corrections. The backreaction of $\\left<T_{\\mu\\nu} \\right>$ on the classical geometry modifies it according to the one-loop corrected Einstein's equation $G_{\\mu\\nu}= 8\\pi G_4\\left<T_{\\mu\\nu}\\right>$. Unfortunately, the stress-energy tensor $\\left<T_{\\mu\\nu} \\right>$ in a black hole spacetime can only be computed approximately, while determining its backreaction is even more difficult \\cite{books}. Only in dimensions $D<4$ was it possible to find exact solutions \\cite{steif,lo,shm,stro}.  On the other hand, an AdS braneworld consists of a bulk AdS$_{D+1}$ space ending on a $D-1$-dimensional domain wall, or brane. A prototype is the RS2 model where AdS$_5$ ends on a $3$-brane, which should model our 3+1 dimensional world. It is therefore natural to look for a suitable description of a black hole in this scenario. However, the attempts to find exact, static, asymptotically flat black hole solutions localized on the brane in AdS$_{D+1>4}$, with regular horizons both on and off the brane, have come empty-handed to date (for published examples see, e.g., \\cite{chr}-\\cite{cfm}). It has even been suggested that static, asymptotically flat, spherical black holes on the brane might not altogether exist in the RS2 model \\cite{bgm}\\footnote{Ref.~\\cite{wiseman} obtains a numerical solution for a static star on an RS2 brane.}. Contrasting this, exact static solutions localized on a 2-brane in AdS$_4$ have been found in \\cite{ehm1,ehm2}.  Here we adopt the point of view that the difficulties in constructing these solutions are no mere accident, but are intricately related to the effects induced by quantum corrections. We use a modification of AdS/CFT correspondence \\cite{malda} for the RS2 model \\cite{apr1}-\\cite{apr7} to connect both problems. Our main result is the following conjecture:  \\begin{quotation} \\noindent {\\it The black hole solutions localized on the brane in the AdS$_{D+1}$ braneworld which are found by solving the classical bulk equations in AdS$_{D+1}$ with the brane boundary conditions, correspond to quantum-corrected black holes in $D$ dimensions, rather than classical ones.} \\end{quotation}  This conjecture follows naturally from the AdS/CFT correspondence adapted to AdS braneworlds. According to it, the {\\it classical} dynamics in the AdS$_{D+1}$ bulk encodes the {\\it quantum} dynamics of the dual $D$-dimensional conformal field theory (CFT), in the planar limit of a large $N$ expansion. Cutting the bulk with a brane introduces a normalizable $D$-dimensional graviton mode \\cite{rs,addk}, while on the dual side this same $D$-dimensional gravity mode is merely added to the CFT, which is also cutoff in the ultraviolet. Then, solving the classical $D+1$-dimensional equations in the bulk is equivalent to solving the $D$-dimensional Einstein equations $G_{\\mu\\nu}= 8\\pi G_D\\left<T_{\\mu\\nu}\\right>_{CFT}$, where the CFT stress-energy tensor incorporates the quantum effects of all planar diagrams. These include particle production in the presence of a black hole, and possibly other vacuum polarization effects.  This conjecture has implications in two directions.  On the one hand, it allows us to view the brane-induced modifications of the metric of a $D$-dimensional black hole as quantum corrections from a CFT, a dual view that sheds light on both problems.  On the other hand, we can use the conjecture to infer, from the known properties of the classical bulk solutions, the properties of the cutoff CFT coupled to gravity. Even if some of the conclusions are derived using the AdS/CFT correspondence, they are typically independent of the existence of a bulk dual: any strongly coupled CFT with a large number of degrees of freedom is likely to behave, when coupled to weak gravity, in a similar manner.  We submit the conjecture to the test by reinterpreting the exact solutions on the 2-brane in an AdS$_4$ braneworld \\cite{ehm1,ehm2} as quantum-corrected, gravitating CFT states in the dual $2+1$ theory, either with or without a negative cosmological constant in 2+1 dimensions, $\\Lambda_3$. As is typical in tests of the AdS/CFT correspondence, the calculations on the CFT side can only be performed at weak 't~Hooft coupling, often at the one-loop order only, and therefore comparisons with the strongly coupled dual of the classical bulk theory, which includes all planar diagrams, are difficult. Even then, we find some instances where the equivalence between the results at weak and strong coupling holds to a great degree of detail.  An interesting spin-off of the analysis is a realization of {\\it quantum censorship of conical singularities}, which we argue is a generic effect independent of the AdS/CFT duality. Gravity in 2+1 dimensions is known to describe massive particles in terms of conical singularities \\cite{djt}. We find that when quantum corrections from a CFT are included, the singularity of a sufficiently massive particle is dressed by a regular horizon. This result is in fact true independently of whether the CFT is strongly or weakly coupled, and acts more efficiently when it has a large number of degrees of freedom.  Since we have a detailed description of the solutions in the AdS$_4$ braneworld, we can apply it to describe the objects which arise in the cutoff CFT. When $\\Lambda_3 = 0$, the theory is characterized by three mass scales: the UV cutoff of the CFT, $\\mu_{UV}$, the 4D Planck mass and the 3D Planck mass, in ascending order. These scales naturally organize the range of CFT configurations into three categories: {\\it (i)} the familiar light CFT states, with masses below the CFT cutoff, which {\\it are not} black holes because of the quantum uncertainty-induced smearing; {\\it (ii)} states with masses between the CFT cutoff and the 4D Planck mass, which also {\\it are not} black holes because of quantum smearing and may receive large quantum corrections in the bulk; and {\\it (iii)} black holes, which are the states with masses above the 4D Planck mass. These black holes may be smaller than the CFT length cutoff, $\\hbar/\\mu_{UV}$, but their description should be reliable since both the bulk and the $2+1$ gravity corrections are small. Our argument that the cutoff CFT can be trusted to distances much shorter than the UV cutoff is analogous to a familiar situation in string theory \\cite{steve}, suggesting that the intermediate mass states and light black holes behave as CFT solitons.  A negative cosmological constant $\\Lambda_3 < 0$, allows for classical BTZ black holes \\cite{btz}. Although the AdS/CFT duality is not fully understood for the case of negatively curved branes, we find that the solutions localized on the 2-brane are naturally interpreted as BTZ black holes with CFT quantum corrections, which are in equilibrium with a thermal bath in AdS$_3$.  There are other localized solutions, all with mass less than $M_{max} = 1/(24G_3)$, with different features, but we find explanations for all of them within the context of our conjecture. Black holes of mass larger than $M_{max}$ are delocalized black strings occupying an infinite region of the bulk, and it is unclear how to describe them within the confines of the $2+1$ theory; in fact, it is likely that such a description should not be possible in terms of only local physics.  In the physically more relevant case of a 3-brane in AdS$_5$ we can not go into a similar level of detail since there are no exact solutions, and classical gravity in $3+1$ dimensions is dynamical. However we can still explore the consequences of our conjecture in a semi-quantitative manner. The description in terms of a CFT coupled to gravity is not reliable until the horizon is larger than the ultraviolet cutoff of the CFT, i.e., the black hole is sufficiently heavy. For these black holes, the CFT+gravity theory allows us to reinterpret the alleged obstruction for finding a static black hole \\cite{bgm} as a manifestation of the backreaction from Hawking effects. The analysis of the trace anomaly of the CFT stress tensor allows us to make this point precise. As long as the anomaly is consistent with the asymptotic AdS$_5$ geometry, the conformal symmetry of the dual CFT is valid in the infrared, and so there is no mass gap. Hence any black hole at a finite temperature will emit CFT modes as a thermal spectrum of Hawking radiation, which on the bulk side is captured by a deformation of the bulk geometry close to the brane, caused by the black hole sourcing the classical gravity equations. We illustrate this to the leading order on the CFT side by showing that the backreaction from Hawking radiation, encoded in the form of a Vaidya-type far-field solution, is consistent with the CFT anomaly. We also discuss the dual bulk picture of Hawking radiation that arises from our conjecture. Within this interpretation, the difficulties encountered in the ongoing quest for the black hole localized on the 3-brane in AdS$_5$ are viewed as a natural, subleading quantum correction to the classical solution, rather than as a no-go theorem for the existence of classical braneworld black holes. ",
        "conclusions": "We have proposed here a radical change of perspective on how to view black holes in the context of AdS/CFT correspondence. The previous work on black holes within the AdS/CFT framework has been aimed at understanding a $D+1$-dimensional black hole sitting at the center of AdS$_{D+1}$ in terms of the quantum states of a CFT at the boundary. In this case, the black hole radiates via quantum effects in the bulk, and one expects to learn about the quantum properties of a black hole by studying its dual boundary description.  Instead, we put the black hole itself in the dual theory extended with dynamical gravity. On the bulk side, this is realized by putting the black hole on a brane in the cutoff AdS bulk, which localizes dynamical gravity. Then we can study the quantum properties of a $D$-dimensional black hole in terms of classical physics in the bulk. The quantum Hawking radiation of CFT modes is described as the emission of gravitational waves into the bulk, and the classical bulk point of view may lead to a better understanding of quantum black hole evaporation. Each of these two approaches prompts different classes of questions, which can be naturally answered within these frameworks.  We have provided strong support for this new point of view with a detailed analysis of the black hole solutions on a 2-brane in AdS$_4$ and their dual $2+1$ CFT+gravity description. Our analysis has also revealed new features of the states of the $2+1$ CFT coupled to $2+1$ gravity, and has shown explicitly that quantum effects can censor singularities. We have found that the main properties of the quantum censorship mechanism in $2+1$ dimensions are in fact quite general, and should remain valid outside of the context of AdS/CFT. The censorship is however amplified in the presence of many CFT modes, and this appears to be the main requirement that makes the quantum censor efficient.  In the context of the RS2 model in AdS$_5$, we have been able to argue why an asymptotically flat, static, regular black hole localized on the brane, could not be found. We emphasize again that while we have been working in the context of AdS braneworlds like RS2, which have proven to be a very useful tool to study black holes, we expect that many of our results should naturally extend to any CFT+gravity theory, even if a dual bulk description along the lines of RS2 does not exist.  There remain a number of open issues. We have given a qualitative argument for why a black hole on a brane should emit classical gravitational waves, but it is still unclear why this emission, which can be analyzed and understood in purely classical terms, should project on the brane as a {\\it thermal} flux of radiation. The problem belongs to a class of connections between classical effects in the bulk and thermal effects in the dual theory. The conventional AdS/CFT approach tried to understand how a state of the CFT encodes the classical causal structure of the bulk black hole. The present problem is quite different and could be an easier one, since we may have some hope of analyzing the classical bulk physics involved in the radiation.  An aspect of our conjecture that we have only barely touched upon is the choice of vacuum of the CFT. This is closely related to understanding Hawking radiation as classical bulk bremsstrahlung. It would be natural to expect that each consistent choice of vacuum should correspond to a specific bulk AdS solution, which differ from each other by the boundary conditions for the bulk waves at the bulk AdS horizons. We have discussed a possible example in the case of BTZ black holes. They admit both the $M=-1/8G_3$ and $M=0$ states as consistent vacua, which we have conjectured to correspond to the two branches of black holes localized on the brane. In $3+1$ dimensions we also had alternative vacua, but we have only examined the physics related to the Unruh vacuum, which models the late time behavior of the collapse. The bulk dual of a black hole with backreaction from the Hartle-Hawking state would be quite interesting as well: The asymptotic thermal radiation is dual to a large black hole inside the AdS$_5$ bulk. The motion of a brane in this spacetime generates the radiation-dominated FRW evolution on the brane. Hence the Hartle-Hawking state should be described in the dual bulk theory as a black hole localized on a brane, which is itself moving in the background of a large bulk black hole in the center of AdS$_5$. The next-to-leading order corrections to the $2+1$ asymptotically flat black holes may lead to a similar picture. On the other hand, the Boulware state should result in a null singularity that is localized on the brane. It would be interesting to check if there exists a relationship between these solutions and the static linearized approximation in the RS2 model \\cite{rs,gartan}. We believe that these questions merit further consideration and hope to return to them in the future."
    },
    "0206/astro-ph0206243_arXiv.txt": {
        "abstract": "Gravitational lensing provides a unique probe of the inner 10--1000 pc of distant galaxies ($z \\sim 0.2$--1). Lens theory predicts that every strong lens system should have a faint image near the center of the lens galaxy, which should be visible in radio lenses but have not been observed. We study these ``core'' images using models derived from the stellar distributions in nearby early-type galaxies. We find that realistic galaxies predict a remarkably wide range of core images, with lensing magnifications spanning some six orders of magnitude. More concentrated galaxies produce fainter core images, although not with any simple, quantitative, model independent relation.  Some real galaxies have diffuse cores and predict bright core images (magnification $\\mm \\gtrsim 0.1$), but more common are galaxies that predict faint core images ($\\mm \\lesssim 0.001$). Thus, stellar mass distributions alone are probably concentrated enough to explain the lack of observed core images, and may require observational sensitivity to improve by an order of magnitude before detections of core images become common. Two-image lenses will tend to have brighter core images than four-image lenses, so they will be the better targets for finding core images and exploiting these tools for studying the central mass distributions of distant galaxies. ",
        "introduction": "Galaxy centers are interesting places to study dynamics and galaxy formation. Their short crossing times make them sensitive to dynamical processes such as relaxation and binary black hole heating (e.g., Ebisuzaki, Makino \\& Okumura 1991; Milosavljevic \\& Merritt 2001). Their deep potential wells collect remnants of the galaxy formation process such as the cores of accreted galaxies (e.g., de Zeeuw \\& Franx 1991; Barnes \\& Hernquist 1992). Their dark matter content provides clues to the interaction between baryons and dark matter by adiabatic compression during galaxy formation (e.g., Blumenthal et al.\\ 1986), and may even reveal properties of the dark matter particle such as cross sections for self-interactions (e.g., Spergel \\& Steinhardt 2000).  Nearby, galaxy centers can be studied directly with high spatial resolution observations. Hubble Space Telescope imaging of early-type galaxies shows that, contrary to theoretical expectations (e.g., Tremaine 1997), the luminosity profiles diverge at small radii (e.g., Faber et al.\\ 1997; Ravindranath et al.\\ 2001; Rest et al.\\ 2001). The profiles seem to fall into two classes: ``core'' galaxies have a distinct transition between a steep outer profile and a shallow inner core that has $I \\propto R^{-\\gamma}$ with $\\gamma \\lesssim 0.3$; while ``power law'' galaxies show no such break and have steep central cusps with $\\gamma \\gtrsim 0.5$. Interestingly, the global properties of the galaxies seem to correlate well with the centers. Core galaxies tend to be luminous, slowly rotating systems with boxy or elliptical isophotes, while power law galaxies tend to be faint, rapidly rotating systems with disky isophotes. Although the division may not be as sharp as originally thought (see Ravindranath et al.\\ 2001; Rest et al.\\ 2001), it still puts strong constraints on the formation process. In hierarchical merging scenarios, simple models cannot easily explain why the large galaxies are so much less dense than their small progenitors (see \\S7 of Faber et al.\\ 1997, and references therein), and some additional process such as heating by binary black holes may be required (e.g., Milosavljevic \\& Merritt 2001; Milosavljevic et al.\\ 2001).  For distant galaxies, we cannot directly resolve 10--100 parsec scales, but we can instead turn to a unique indirect probe offered by gravitational lensing. Lens theory predicts that if the central mass distribution is shallower than $\\rho \\propto r^{-2}$ then any multiply-imaged gravitational lens must have an odd number of images (Burke 1981; Schneider, Ehlers \\& Falco 1992). Standard image configurations\\footnote{The rare exception is a configuration where the source lies in a naked cusp (e.g., Schneider et al.\\ 1992); among more than 60 known lenses, APM~08279+5255 is the only candidate naked cusp lens (Lewis et al.\\ 2002). The only other exception is B1359+154, a unique lens where three lens galaxies jointly produce six bright images, and models predict three additional core images (Rusin et al.\\ 2001).} have two or four bright images lying $\\sim$3--10 kpc from the center of the lens galaxy, with the remaining image just 10--100 pc from the center and much fainter than the others. Because the ``core'' image is very sensitive to the central surface density of the lens galaxy, with a higher density corresponding to a fainter image, it offers a unique way to constrain the density on scales that cannot be directly resolved. This probe of galaxy centers can in principle be applied to all lens galaxies, which now number more than 60 and are predominantly early-type galaxies spanning the redshift range $z \\sim 0.3$--1 (e.g., Kochanek et al.\\ 2000). It is conceptually equivalent to using radial arcs to constrain the central profiles of lensing clusters (e.g., Mellier, Fort \\& Kneib 1993; Smail et al.\\ 1996; Molikawa \\& Hattori 2001; Oguri, Taruya \\& Suto 2001).  The best observational data on core images come from radio lenses, because the lack of radio emission from most lens galaxies enables sensitive searches for core images. The Cosmic Lens All-Sky Survey found 18 radio lenses but no core images, based on radio maps where the dynamic range is typically several tens to several hundreds but reaches 1200 for B1030+074 and 2000 for B0218+357 (Rusin \\& Ma 2001; Norbury et al.\\ 2002). Several other radio lenses have candidate core images (MG~1131+0456, Chen \\& Hewitt 1993; PMN J1632-0033, Winn et al.\\ 2002) although the hypothesis that the central radio flux originates in the lens galaxy cannot be ruled out. At optical and near-infrared wavelengths, APM~08279+5255 has an odd number of images (Ibata et al.\\ 1999; Egami et al.\\ 2000), but its interpretation is not clear. The third image may be a core image, in which case it indicates a large low-density core in the lens galaxy (Ibata et al.\\ 1999; Egami et al.\\ 2000; Mu\\~noz, Kochanek \\& Keeton 2001), or it may be a case of a ``naked cusp'' image configuration, in which case it contains no information about the center of the lens galaxy (Lewis et al.\\ 2002). No other optical core images have been seen, although the searches are of course hindered by light from the lens galaxies.  The apparent discrepancy between data and theory provides the desired opportunity to learn about the centers of distant galaxies. Motivated both by theoretical expectations (see Tremaine 1997) and by ease of use, many analyses have assumed models with a finite density core and obtained limits on lens galaxy core radii (e.g., Narayan, Blandford \\& Natyananda 1984; Narasimha, Subramanian \\& Chitre 1986; Blandford \\& Kochanek 1987; Hinshaw \\& Krauss 1987; Narayan \\& Schneider 1990; Wallington \\& Narayan 1993; Kochanek 1996; Evans \\& Hunter 2002). Rusin \\& Ma (2001) instead used power law models and obtained a lower limit on the power law index, $\\gamma > 0.8$ at 90\\% confidence for a surface density $\\Sigma \\propto R^{-\\gamma}$. The question remained, though, whether these two classes of models were realistic enough to provide robust, model independent conclusions about lens galaxy centers. Mu\\~noz et al.\\ (2001) introduced double power law models where the core region is allowed to have a power law cusp whose index is independent of the density profile at large radii. They found that the lack of a core image in B1933+503 robustly implies $\\gamma \\gtrsim 0.6$ for that one galaxy. Keeton (2001) studied core images statistically using models with cuspy stellar components, treated as generalized Hernquist (1990) models, embedded in dark matter halos. He found that the models were inconsistent with the data, perhaps because generalized Hernquist models may not accurately represent the stellar components of galaxies on 10--100 pc scales.  The goal of this paper is to reconsider the core image problem using more realistic models derived from nearby galaxies, and more generally to discuss using core images as tools for studying the centers of distant ($z \\sim 0.3$--1) galaxies. Nearby early-type galaxies have surface brightnesses that can be modeled as a Nuker law (Lauer et al.\\ 1996; Byun et al.\\ 1996), and in \\S2 we discuss lensing by such galaxies. In \\S3 we consider in a general way what physical properties of lens galaxies determine core images properties, or conversely what we can learn about lens galaxies by studying core images. In \\S4 we study in detail the core images predicted by a sample of realistic lens galaxies. Finally, in \\S5 we offer a summary and conclusions.  We assume the popular $\\Lambda$CDM cosmology with matter density $\\Omega_M=0.3$, cosmological constant $\\Omega_\\Lambda=0.7$, and Hubble constant $H_0=75$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "Core images in strong gravitational lens systems provide a unique probe of the centers of galaxies at redshifts $z \\sim 0.2$--1. The brightnesses of core images are determined by the density profiles of galaxies inside $\\lesssim$200 pc. The lack of core images in observed lenses, especially in radio lenses, sets strong lower limits on the central densities of the lens galaxies (e.g., Wallington \\& Narayan 1993; Mu\\~noz et al.\\ 2001).  The mapping between core images and galaxy centers can be studied in two directions. In the forward problem knowledge of galaxy centers is used to make predictions about core images. Based on the first lens models drawn directly from the resolved stellar mass distributions of nearby early-type galaxies, we predict that real galaxies should produce a remarkably wide range of core images. Some should have bright core images (magnification $\\mm \\gtrsim 0.1$), while many others will have core images that are faint ($\\mm \\lesssim 0.001$) or absent altogether.  Qualitatively, more concentrated galaxies produce fainter core images. Quantitatively, however, there does not seem to be a simple predictive relation between observed galaxy properties and core images. Lensing is biased against galaxies with bright core images, because they have smaller cross sections than comparable galaxies with faint core images. Four-image lenses should tend to have fainter core images than two-image lenses, because in quads the source is always close to the center of the lens galaxy where the core image magnification is low. Supermassive black holes in the centers of galaxies can suppress faint core images, but they have little effect on bright core images or on the mean magnification.  The connection between core images and galaxy centers can also be studied in the inverse problem, where the analysis of core image data yields constraints on the centers of lens galaxies. Previous work placed limits on the core radius or on the logarithmic slope of the density (e.g., Wallington \\& Narayan 1993; Rusin \\& Ma 2001), but it was not clear how model dependent those constraints were. We obtain a general statement of the connection between the density profile and core images: the mean core image magnification $\\ma$ is inversely related to the density at the lensing critical radius $R_{\\rm rad}$ (eq.~\\ref{eq:mu5}); and this critical radius is determined by the shape of the density profile (eq.~\\ref{eq:k2}), with more concentrated galaxies corresponding to smaller critical radii and fainter core images. Unfortunately, neither our general formalism nor our Nuker law lenses suggest any simple, model independent measure of the mass concentration that determines the core image properties. The interpretation of core image data will therefore continue to rely on detailed models of individual lenses. The model dependence can be held in check, though, by using general models like the Nuker law or the cuspy lenses introduced by Mu\\~noz et al.\\ (2001), as opposed to overly simple flat core or pure power law models.  We conclude that in many cases the stellar mass in lens galaxies is probably concentrated enough to render core images faint (also see Evans \\& Hunter 2002). This is not to say that bright core images cannot exist --- certainly there are realistic galaxies that predict bright core images, and the probability that they are selected for lensing is non-zero. But the fact that core images have not yet been found (with perhaps one or two exceptions) is probably not a surprise. As the search continues, two-image lenses should be better targets than four-image lenses for revealing core images."
    },
    "0206/astro-ph0206075_arXiv.txt": {
        "abstract": "{We have performed new determinations of sulphur and silicon abundances for a sample of 26 disk stars based on high-resolution, high signal-to-noise spectra. The results indicate a solar $\\SFe$ for $\\feh >-0.3$, below which $\\SFe$ increases to $\\sim$0.25 dex at $\\feh =-1.0$. We find that there is a good correlation between $\\sh$ and $\\sih$, indicating the same nucleosynthetic origin of the two elements. It seems that the ratio of sulphur to silicon does not depend on metallicity for $\\feh > -1.0$. The implications of these results on models for the nucleosynthesis of $\\alpha$-capture elements and the chemical evolution of the Galaxy are discussed. ",
        "introduction": "The trend of $\\OFe$ vs. $\\feh$ has recently been paid much attention because works by different authors yield very dissimilar results. Specifically, Israelian et al. (\\cite{Israelian98}, \\cite{Israelian01a}) and Boesgaard (\\cite{Boesgaard99}) derived oxygen abundances from OH UV lines and the infrared $\\Oone$ triplet and found a linear rise of $\\OFe$ to  1.0 dex when the metallicity decreases to $\\feh =-3.0$ dex. Some doubts have, however, recently been cast on the correctness of these results from calculations based on the new generation of 3D hydrodynamical model atmospheres (Asplund \\& Garc\\'{\\i}a P{\\'e}rez \\cite{Asplund01}). Other studies (e.g. Nissen et al. \\cite{Nissen92}, Nissen et al. \\cite{Nissen01}) based on the forbidden line at $\\lambda$ 630.03\\,nm suggest, on the other hand, a more or less flat trend of $\\OFe$ in the metallicity range of $-2.0 < \\feh < -1.0$. Concerning other well-studied $\\alpha$-capture elements, e.g. Mg, Si and Ca, earlier observed trends of $\\afe$ versus $\\feh$ (e.g. Ryan et al. \\cite{Ryan96}) seemed to be in  good agreement with a plateau of $\\afe$ $\\sim$ 0.4 dex. In particular, Fuhrmann (\\cite{Fuhrmann98}) found a clearly flat $\\MgFe$ ratio for $\\feh < -0.8$ based on a LTE abundance analysis of Mg and Fe. However, the recent work of Stephens \\& Boesgaard (\\cite{Stephens02}) based on Keck spectra of more than 50 halo and disk stars suggests a quasi-linear trend of increasing $\\afe$ with decreasing $\\feh$ (see their Fig. 20), but we note that their average $\\afe$ is only about $\\sim$ 0.4 dex at $\\feh = -3.0$. In addition, Idiart \\& Thevenin (\\cite{Idiart00}) have found a complicated structure in $\\MgFe$ and $\\CaFe$ vs $\\feh$ based on a non-LTE re-analysis of equivalent widths from various sources. Although the scatter is significant for $\\feh < -1.0$, their results are consistent with plateau-like mean trends at $\\sim +0.2-0.3$ for both $\\MgFe$ and $\\CaFe$ for halo stars. Since different trends of $\\afe$ vs. $\\feh$ correspond to different nucleosynthesis histories of the $\\alpha$ elements, a clarification of the divergency on the trend of $\\OFe$ vs. $\\feh$ is very important for many astrophysical fields including stellar evolution theory, nucleosynthesis theory, the formation and chemical evolution of the Galaxy, and the age of the Galaxy and the Universe.  Another $\\alpha$ element, sulphur, is highly interesting as an independent tracer of the nucleosynthesis of $\\alpha$ elements and a well-established trend of $\\SFe$ vs. $\\feh$ could be helpful in clarifing the discrepancy of the $\\OFe$ ratio. Furthermore, sulphur is of particular interest in connection with studies of the chemical enrichment of damped Ly$\\alpha$ systems being one of the few elements that is not depleted onto dust (Centuri\\'{o}n et al. \\cite{Centurion00}). Unfortunately, sulphur has been ignored for many years probably due to the lack of measurable lines in the optical spectra. Presently, our knowledge of sulphur abundances is limited to only a few studies: Clegg et al. (\\cite{Clegg81}), Fran\\c{c}ois (\\cite{Francois87}, \\cite{Francois88}) and recent works by Israelian \\& Rebolo (\\cite{Israelian01}) and Takada-Hidai et al. (\\cite{Takada02}). Again, the disagreement on the $\\SFe$ ratio at low metallicities appears: Fran\\c{c}ois (\\cite{Francois87}, \\cite{Francois88}) suggested the $\\SFe$ ratio to be constant for stars with $-1.6 < \\feh <-1.0$, while Israelian \\& Rebolo (\\cite{Israelian01}) favor a linear rise of $\\SFe$ ratio with decreasing $\\feh$ in the metallicity range of $-3.0 < \\feh <-0.6$. With a smaller slope ($-$0.25 vs $-$0.46 in Israelian \\& Rebolo paper) the work of Takada-Hidai et al. (\\cite{Takada02}) support a linear trend of $\\SFe$.  Sulphur and iron abundances in disk stars are usually better determined than in the case of halo stars. The sulphur lines have a suitable strength for accurate abundance determination and $\\Fetwo$ lines, which have very small non-LTE effects, can be used for deriving the iron abundance. Therefore, the $\\SFe$ vs. $\\feh$ trend can be more reliably established. With this advantage, the determination of sulphur abundances in disk stars with $\\feh > -1.0$ is important especially from the nucleosynthesis theory point of view. The comparison of the sulphur behaviour with those of other $\\alpha$ elements in disk stars provides a key clue to the genuine origin of sulphur and further to the chemical evolution of the Galaxy. As the nearest even-$Z$ neighbour of sulphur in the periodic table of elements, silicon is very suitable as a representative of the $\\alpha$ elements for such a comparison.  In the present work, we investigate the behaviour of sulphur in disk stars and the nucleosynthesis origin of this element. The investigation takes benefit of more accurate observations of sulphur lines, better determinations of stellar parameters, improved atmospheric models, and hence more reliable derivations of sulphur, silicon and iron abundances than in previous works (Clegg et al. \\cite{Clegg81}; Fran\\c{c}ois \\cite{Francois87}, \\cite{Francois88}), where the scatter of $\\SFe$ for disk stars is so large that it is not clear if sulphur really behaves as an $\\alpha$-capture elements or not. In a following paper (Nissen et al. \\cite{Nissen02}), the study will be extented to halo stars such that the trend of $\\SFe$ will be covered for the whole metallicity range of  $-3.0 < \\feh < +0.2$. ",
        "conclusions": "We have carried out new measurements of sulphur and silicon abundances for 26 main sequence stars in the metallicity range of $-1.0 < \\feh < +0.5$ and have established the trends of $\\SiFe$ and $\\SFe$ with a small scatter at a given metallicity. The result indicates a very good correlation between sulphur and silicon and the correlation does not depend on the metallicity. The results provide support for the chemical evolution models of Goswami \\& Prantzos (\\cite{Goswami00}), Timmes et al. (\\cite{Timmes95}), Chiappini et al. (\\cite{Chiappini97}) and Samland et al. (\\cite{Samland98}) for these two elements in the case of disk stars. It is concluded that the nucleosynthesis of sulphur is a proxy for silicon, which is the nearest $\\alpha$-capture element of sulphur in the periodic table of elements."
    },
    "0206/astro-ph0206305_arXiv.txt": {
        "abstract": "We analyze the power spectra of the lightcurves of long gamma-ray bursts, dividing the sample in bins of luminosity, using the recently discovered variability-luminosity correlation. We find that the value of the variability parameter strongly correlates with the frequency that contains most of the power in the burst comoving frame. We compute the average power spectra in luminosity bins. The average power spectrum is well described by a broken power-low and the break frequency is a function of the variability parameter, while the two slopes are roughly constant. This allow us to conclude that scattering processes do not play a relevant role in modelling the lightcurves. We finally discuss in which conditions scattering may still play a relevant role in shaping the spectra of GRBs. ",
        "introduction": "The study of the BATSE $\\gamma$-ray bursts (GRBs) lightcurves has recently gained new interest thanks to the discovery of the variability-luminosity (Fenimore \\& Ramirez-Ruiz 2000; Reichart et al. 2001) and lag-luminosity (Norris et al. 2000) correlations (see also Chang, Yoon \\& Choi 2002). These enable to assign a tentative redshift to BATSE GRBs, allowing to perform spectral and temporal analysis of the lightcurves in the burst comoving frame, where their properties are more closely linked to the physics of the burst itself. As an example, the use of these correlations enabled the possible discoveries of an evolution of the luminosity function with redshift (Lloyd-Ronning, Fryer \\& Ramirez-Ruiz, 2002) and of a correlation between the peak photon frequency with the luminosity of the GRB (Lloyd-Ronning \\& Ramirez-Ruiz 2002).  The variability-luminosity correlation predicts that the peak luminosity of a burst is correlated to its degree of variability, whose operational definition is related to the normalized variance, or the root mean square of the deviations from a smoothed version of the light curve. More variable lightcurves have larger luminosities. The lag-luminosity correlation is instead based on the measure of a temporal lag between the detection of high energy and low energy photons. It is found that the more the burst is luminous, the smaller is the lag. Besides being extremely useful tools, these correlations also call for an explanation of their origin. Several possible interpretations have been discussed in the literature. In particular, it is shown that an underlying correlation between the isotropic equivalent luminosity and the Lorentz factor of the fireball can explain both the correlations (Ramirez-Ruiz \\& Lloyd-Ronning 2002; Kobayashi, Ryde \\& MacFadyen 2002; Salmonson 2000; \\mes~et al. 2002). In some of these works, however, a significant role is attributed to the modification of the temporal properties of the light curves as a consequence of scattering by cold or hot electrons (see also Panaitescu, Spada \\& \\mes~1999; Spada, Panaitescu \\& \\mes~2000). The role of scattering seem to be strengthened by the need of explaining the peak photon frequency vs. isotropic equivalent luminosity correlation (Lloyd-Ronning \\& Ramirez-Ruiz 2002), which is not naturally predicted in the internal shock scenario (Ghisellini, Celotti \\& Lazzati 1999).  In this paper, we concentrate on the variability luminosity correlation, using power spectra as a diagnostic to investigate several related issues. First (\\S 2), we compute the dominant frequency of the power spectrum for each GRB in our sample of 220 lightcurves. We find that this dominant frequency strongly correlates with the variability measure. We then divide the sample in bins of variability and compute the average power spectrum. We find (\\S 3) that the power spectra are self-similar (Beloborodov, Stern \\& Svensson 1998; 2000; hereafter B98 and B00; Chang \\& Yi 2000) in all the variability bins, with a break at a frequency that correlates with the variability parameter. No sign of a cut-off in the spectrum at large frequencies is observed. We (\\S 4) develop a shot-noise model for the lightcurve, taking into account the effect of scattering in Fourier space (\\S~5). In \\S 6 we discuss our results, showing how scattering processes cannot be responsible for the variability-luminosity correlation, and also constraining the regions of the parameter space where photon scattering on cold electrons may imprint detectable signatures on the photon spectra of GRBs without leaving a detectable trace in their power spectra. ",
        "conclusions": "How do the shot-noise theory compares to the average spectra derived in \\S~3? The observed spectra show a clear break, which we show is strongly correlated with the variability parameter. The power-law slope changes from $-2/3$ before the break to $\\sim -2$ after the break.  Let us for the moment neglect photon scattering. The fact that the largest slope is $\\sim-2$ suggests that an exponential shot noise model can easily reproduce the observations, with a typical pulse decay time longer than the rise time (Norris et al. 1996). The fact that the low frequency slope is not flat, can be easily accommodated invoking a power-law distribution of pulse durations (cfr. Eq.~\\ref{eq:alp}) \\begin{equation} n(\\tau_d)\\propto\\tau_d^{-1/3} \\label{eq:sixt} \\end{equation} The correlation between the break frequency and the variability (or the GRB luminosity) can therefore be interpreted as a correlation between the typical shortest relevant pulse duration and the luminosity.  This does not mean that the lightcurve cannot have shorter pulses. In fact, if the distribution of pulse durations has a break, shorter pulses can be present but not contribute to the PSD.  Let us now consider the effect of photon scattering. In a very simple scenario, all the pulses may go through a scattering screen with opacity $\\tau$ and radius $R$. In this case, a clear break should be present in the PSD at a frequency $\\nu=c/(\\pi\\,R\\,\\tau)$. It is straightforward to show that the observed break cannot be due to scattering but must, instead, be attributed to the intrinsic properties of the unscattered pulses. In fact, the change of slope before and after the break is only of unity, while scattering would require either a jump of 2 in slope (intermediate opacity) or an exponential cutoff ($\\tau>10$). Alternatively, one can consider a scenario in which the opacity and size of the screen are different for different pulses. In order to be consistent with the lack of a pronounced break in the observations, this requires that the scattering time scale $\\tau R/c$ is smaller than $\\tau_d$. In this case, the PSD of the single pulse would be affected only in the power-law tail and its effect would not be observed in the average PSD. In order to explain the break in the average PSD, however, one should consider a broken power-law of pulse durations. With such conditions one can explain the observed average power-spectrum. However, the role of scattering is not relevant in the shape of the PSD, and so it cannot be important in the measure of the variability parameter $V$. It may be also envisaged a case in which $\\tau R/c\\gg\\tau_d$ for all pulses, with a broken power-law distribution of the values of $\\tau R$ suited to mimic the average PSD shape. In this case, however, the lightcurve would be entirely dominated by scattering. The pulses should have the shape of the transfer functions $k(t)$, in contradiction with what is found in their direct analysis (Norris et al. 1996).  A more detailed modelling of an internal shock process should however take into account that different pulses may not be scattered with the same value of $\\tau R/c$ since the shortest pulses are likely to be produced closer to the engine, where the relativistic wind is more dense and opaque. To mimic such a case, we consider a lightcurve in which only the shortest pulses underwent scattering. In this case, the PSD should show a steepening break followed by a flattening break when the unscattered pulse component becomes dominant (see lowest right panel of Fig.~\\ref{fig:sim}). Again, this is not observed in real data (see Fig.~\\ref{fig:ave}).  It may be argued that these models are too simplistic, since in real simulations each pulse is smeared with its own value of the parameter $\\tau{}R/c$. What emerges from our analysis is that no clear signature of scattering is present in the PSD data.  Is may still be possible that a more detailed numerical simulation for the evolution of the flow can include scattering in such a way that its effect is relevant but does not produce a cutoff in the PSD. This requires however a fine tuning of the distribution of the parameter $\\tau{}R/c$ and that the shape of at least part of the pulses is dominated by the transfer functions of Fig.~\\ref{fig:ker}. Such simulations, with a proper treatment of the scattering, are called for if the proposed link between scattering and variability has to be believed.  Even though scattering processes are shown to be not important in determining the variability parameter $V$, it is still possible that they have some relevance in shaping the spectra of GRBs. In fact, if the scattering screen is small and thick, the cutoff frequency can be large, in a range in which the PSD is dominated by the white noise. This kind of scattering would not be influent for the temporal pulse profile, but the photon energy may be shifted by \\begin{equation} \\Delta\\epsilon\\simeq{{\\tau^2\\epsilon}\\over{m_e\\,c^2}}\\left(4kT-\\epsilon\\right) \\end{equation} where $T$ is the temperature of the scattering electrons. This mechanism for modifying the typical energy of the photons in GRB spectra has been proposed and discussed, e.g., in Ramirez-Ruiz \\& Lloyd-Ronning (2002) and \\mes~et al. (2002). In particular, assuming that the comoving electron temperature is small and asking that the photon energy is sizably modified ($\\Delta\\epsilon\\sim\\epsilon$) one can show that the opacity must be \\begin{equation} \\tau\\sim\\sqrt{{{m_e\\,c^2}\\over{\\epsilon}}} \\approx \\sqrt{\\Gamma} \\end{equation} where $\\Gamma$ is the bulk Lorentz factor of the flow and the rightmost term holds if the initial observed energy of the photon was closed to $m_e\\,c^2$.  Before discussing this issue in detail, we consider that all the computation described above are relevant if the scattering screen is comoving with the relativistic flow. In fact, in order to preserve the burst variability, a scattering screen at rest in the frame of the host galaxy needs to have an optical depth much smaller than unity. For this reason, the screens that we consider here are comoving with the flow, they can be the same shell in which the radiation is produced, or the total contribution of previously ejected shells.  In order for the scattering to be important in shaping the spectrum of GRBs, one have to assume a large opacity, so that the cutoff in the power spectrum should be exponential. Consider the bursts with largest variability (lower left panel of Fig.~\\ref{fig:ave}). The spectrum is well described by a power-law up to observed comoving frequencies $\\nu\\sim10$~Hz. We can then constrain the comoving size of the scattering screen to be \\begin{equation} R^\\prime < 10^9 \\,\\Gamma^{1/2} \\label{eq:con} \\end{equation}  In the framework of internal shock, the comoving width of the shell is given by $\\Delta^\\prime=r/\\Gamma\\sim r_0\\Gamma$, where $r_0$ is the size of the shell at the moment of the ejection. The constraint~\\ref{eq:con} implies then \\begin{equation} r_0\\lsim10^8\\,\\Gamma_2^{-1/2} \\end{equation} which is consistent with the internal shock picture only if the inner engine produces many shells, of the order of 1000 shells in a burst lasting for 10 seconds. Without being confined in the standard internal shock model, one can envisage a scenario in which the energy is liberated in the fireball at radii $r<r_0 \\Gamma^2$. In this case, it must be considered that the photons will be advected by the flow, which will become optically thin to radiation (due to the shell expansion) in a time scale smaller than the diffusion time scale of the photons. In this case, the break frequency would be expected to be at $\\nu\\simeq c\\Gamma^2/(2\\pi\\,R)$, which is always above the measured 10~Hz."
    },
    "0206/astro-ph0206419_arXiv.txt": {
        "abstract": "{ We present deep optical imaging of three face-on disk galaxies together with a detailed description of the reduction and calibration methods used, in order to measure the intrinsic shape of their outer stellar edges. Whereas it is now well accepted that disks of spiral galaxies are not infinite exponential beyond galactocentric distances of about 3-5 radial scalelengths, the genuine structure of the truncation region is not yet well known. Our data quantitatively establish a smooth truncation behaviour of the radial surface brightness profiles and is best described by a two-slope model, characterised by an inner and outer exponential scalelength separated at a relatively well defined break radius. This result disagrees with the frequently assumed sharply truncated nature of the radial surface brightness profiles and implies the presence of stars and even star-formation beyond the break radius. In addition, we do not find a strong influence of a nearby companion on the ratio of the break radius to the radial scalelength. Our results denote new observational constraints for the search of the physical explanation for these smooth disk truncations. ",
        "introduction": "Since the pioneering work of \\cite{devauc1,devauc2} and \\cite{free} the radial surface brightness, tracing the stellar distribution, is empirically known to be well described by an exponential decline. The inner part can be dominated by a bulge, bar, or ring component, whereas the outer disk follows an exponential quite well with only minor deviations related to spiral arms. However, \\cite{vdk79} has already pointed out that the profiles appear to break off towards the outer parts. This is later quantitatively confirmed using optical filters by \\eg \\cite{vdk82a}, \\cite{bd}, \\cite{pohlen1,pohlen2}, and \\cite{dgkrwe}. Just recently, data by \\cite{floridoNIR} for the first time suggest these truncations to be also present in NIR images. These so called cut-off radii are practically only observed in edge-on galaxies due to the preferred line-of-sight integration. While the existence of the truncations is now well accepted there are still two major unsolved problems.  \\newline On the one hand we are still missing a unique physical interpretation to describe this observational phenomenon. The proposed hypotheses span a rather wide range of possibilities. \\Citet{vdk87} has deduced a connection to the galaxy formation process describing the truncations as remnants from the early collapse. \\cite{ferg3}, for example, have proposed an evolutionary scenario represented by the viscous disk evolution models. \\cite{kenni} has suggested a dynamical critical star-formation threshold, and \\cite{batta} have proposed a connection with large scale galactic magnetic fields. \\newline On the other hand the detailed structure of the radial surface brightness distribution of disk galaxies is still not well constrained, although already  \\cite{naeslund}, \\cite{byun}, \\cite{dgkrwe}, and \\cite{pohlenphd} have found that the truncations are not that sharp as thought before. While the former studies assume the truncation region to be exponential, \\cite{floridoNIR} used a model independent deprojection method to measure and analyse the truncation curve (difference between measured profile and an infinitely exponential decline). However, nearly all studies used edge-on galaxies to search for truncations, whereas only the profiles of face-on galaxies are practically independent of possible line-of-sight effects caused by the integration across a multicomponent disk and less affected by the disturbing dust influence. Observing face-on galaxies is therefore mandatory to study the detailed quantitative behaviour at the truncation region and to determine the scalelengths in an unbiased way.  Up to now \\cite{vdkshos82} and \\cite{shos} have reported the only known quantitative values of cut-offs in the literature for face-on galaxies (NGC\\,3938 and NGC\\,628). \\Citet{vdk82a} have just found indications of a radial truncation for NGC\\,5033 and NGC\\,3198 and no evidence in NGC\\,4258 and NGC\\,5055, all being late type galaxies. Although calling the observed feature near the edge of the disk a 'truncation' \\cite{vdk82a} have already stated that ``the edges are not infinitely sharp (we usually derive an upper limit to the scalelength beyond $R_{\\rm max}$ of about 1\\,kpc)''. In a subsequent paper \\cite{vdk88} has stated that out of the 20 face-on galaxies observed by \\cite{wev} only four (NGC\\,2681, NGC\\,4203, NGC\\,5005, NGC\\,6340) ---all being early type galaxies and therefore probably possessing a different disk formation history--- do not show any sign for a drop off as judged from the last three contours. He has also recognized that only NGC\\,5371 shows a relatively sharp edge but without quoting detailed numbers. Explaining this in addition to the much lower surface brightness for face-on systems, with the problem of dealing with intrinsic deviations from the circular symmetry of disks which could be hidden by an azimuthally averaged profile. These asymmetries, such as spiral arms, especially from the young population should unavoidably contribute in less inclined systems. All these results are based on photographic material and up to now there is no systematic follow up of truncations for face-on galaxies. Especially a detailed quantitative modelling of radial surface brightness profiles based on higher quality CCD imaging data is missing.  However, there are at least two recent papers in the literature which both have not found any sign for a truncation investigating two nearly face-on galaxies. \\cite{barton} ---with a sensitivity limit of $\\mu^{\\rm V}_{\\rm lim}\\eq29.9$\\,\\magsqarcsec--- have not found any evidence for a cut-off for the galaxy NGC\\,5383 out to approximately three Holmberg radii. They have measured an exponentially falling disk with several distinct departures from the smooth decline (two spiral arms with one outside the Holmberg radius), but their fitting yields unsatisfactory results. Nevertheless, in addition to a nearby companion (UGC\\,8877) this galaxy unfortunately shows a rather complex morphological structure (type: PSBT3*P according to the RC3\\footnote{Third Reference Catalogue of bright galaxies \\citep[RC3:][]{rc3}.}) and has also been rejected by \\cite{vdk88} in his study. \\cite{weiner} also have found no sign of a truncation out to ten scalelength for the NGC\\,4123 (type: SBR5 + starburst according to NED\\footnote{NASA/IPAC Extragalactic Database (NED).}) in their azimuthally averaged light profile. These studies have in common that both target galaxies are strongly barred and show a rather peculiar overall morphology.  The detailed structure of the truncation region can be expected to give important constraints for the physical nature of these outer stellar edges. In addition our analysis has consequential effects on the determination of the truncation as a structural parameter which is important to compare observational results with model predictions. Depending on the true physical nature these outer edges may contain fossil evidences imprinted by the detailed galaxy formation and evolutionary history. ",
        "conclusions": "Our deep optical imaging of three face-on galaxies confirms the finding of \\cite{naeslund}, \\cite{byun}, \\cite{dgkrwe}, and \\cite{pohlenphd} that the observed truncations of galactic stellar disks are not infinitesimally sharp but better described by an additional exponential decline. Since face-on scalelengths are not influenced by line-of-sight integration problems as in the case for edge-on galaxies (so far investigated), we are able to quantitatively introduce a new two-slope model. We find that the radial surface brightness profiles are best represented by a two-slope exponential profile, characterised by an inner \\hin and outer \\hout scalelength separated at a break radius \\rbr, rather than the previously assumed sharply-truncated exponential model proposed by \\cite{vdk81a} and used for example in the study of \\cite{pohlen2}. The mean value for the distance independent ratio of break radius to inner scalelength which marks the start of the truncation region is $\\rbrdhin\\!=\\!3.9\\pm0.7$.  Although larger statistics are also definitely needed, the lack of any conspicuity in the parameters for NGC\\,5434 with its close physical companion points against an interaction hypothesis for explaining stellar truncations. In addition, we could not correlate undoubtedly the end of the visual spiral arms with the occurrence of the break radii, which may be a hint relating the truncation phenomenon to star formation properties. Similar deep imaging of grand-design spirals are needed to clarify this hypothesis. Since we do not have deep \\halpha maps for our galaxies we have searched the literature for signs of break radii in published surface brightness profiles and have compared them in the case of NGC\\,628 with already available \\halpha\\,data. Although not conclusive in this case an \\halpha\\,follow up of the three observed face-on galaxies will be a promising project to determine the origin for the observed break radii (or smooth truncations) and address the star formation threshold hypothesis by \\cite{kenni}.  However, the detailed quantitative structure of the {\\it smoothness} of the observed truncations presented here will set strong constraints for present and future theoretical models trying to explain this observational phenomenon."
    },
    "0206/astro-ph0206133_arXiv.txt": {
        "abstract": "The discovery of the second highly $r$-process-enhanced, extremely metal-poor star, CS~31082-001 ([Fe/H] $= -2.9$) has provided a powerful new tool for age determination, by virtue of the detection and measurement of the radioactive species uranium and thorium.  Because the half-life of $^{238}$U is one-third that of $^{232}$Th, the U-Th pair can, in principle, provide a far more precise cosmochronometer than the Th-Eu pair that has been used in previous investigations.  In the application of this chronometer, the age of (the progenitor of) CS~31082-001 can be regarded as the minimum age of the Galaxy, and hence of the universe.  One of the serious limitations of this approach, however, is that predictions of the production ratio of U and Th have not been made in the context of a realistic astrophysical model of the $r$-process. We have endeavored to produce such a model, based on the ``neutrino winds'' that are expected to arise from the nascent neutron star of a core-collapse supernova.  In this model, the proto-neutron star mass and the (asymptotic) neutrino sphere radius are assumed to be $2.0 M_\\odot$ and 10~km, respectively.  Recent hydrodynamic studies indicate that there may exist difficulties in obtaining such a compact (massive and/or small in radius) remnant.  Nevertheless, we utilize this set of parameter choices since previous work suggests that the third $r$-process peak (and thus U and Th) is hardly reached when one adopts a less compact proto-neutron star in the framework of the neutrino-wind scenario.  The temperature and density histories of the material involved in the neutron-capture processes are obtained with the assumption of a steady flow of the neutrino-powered winds, with general relativistic effects taken into account. The electron fraction is taken to be a free parameter, constant with time. The $r$-process nucleosynthesis in these trajectories is calculated with a nuclear reaction network code including actinides up to $Z = 100$.  The mass-integrated $r$-process yields, obtained by assuming a simple time evolution of the neutrino luminosity, are compared to the available spectroscopic elemental abundance data of CS~31082-001. As a result, the ``age'' of this star is determined to be $14.1 \\pm 2.5$~Gyr, in excellent agreement with lower limits on the age of the universe estimated by other dating techniques, as well as with other stellar radioactive age estimates. Future measurements of Pt and Pb in this star, as well as expansion of searches for additional $r$-process-enhanced, metal-poor stars (especially those in which both U and Th are measurable), are of special importance to constrain the current astrophysical models for the $r$-process. ",
        "introduction": "Uranium and thorium are regarded as potentially useful cosmochronometers because their long radioactive decay half-lives ($^{232}$Th: 14.05 Gyr, $^{238}$U: 4.468 Gyr) are significant fractions of the expected age of the universe.  All of the actinides, including uranium and thorium, are thought to be pure $r$-process elements, since the alternative neutron-capture process, the $s$-process, terminates at $^{209}$Bi. The excellent agreement of the relative abundances of neutron-capture elements in the extremely metal-poor (${\\rm [Fe/H]} = - 3.1$), $r$-process-element-enhanced star CS~22892-052 with the solar $r$-process pattern initially suggested that thorium might serve as a precise cosmochronometer \\citep{Sned96}. The time that has passed since the production of the thorium that is now observed in the outer atmosphere of such an old halo star can be regarded as a hard lower limit on the age of the universe.\\footnote{For convenience of nomenclature, in the context of radioactive cosmochronology it is common practice to refer to the ``age'' of a given star as the time difference between the production event (or events) of the radioactive (and stable $r$-process) species that have been incorporated into the outer atmosphere of a low-mass ($< 0.8 M_\\odot$) star that has not yet evolved past the giant branch stage of evolution, and hence is amenable to spectroscopic study.}  Sneden et al. (1996) estimated the age of this star to be $15.2 \\pm 3.7$~Gyr, comparing the measured Th/Eu ratio with an initial production ratio constrained by theory. This age limit, albeit with a rather large error bar, is in good agreement with the age limits derived by completely different techniques, e.g., for globular clusters ($12.9 \\pm 2.9$~Gyr; Carretta et al. 2000), for Type Ia supernovae ($14.9 \\pm 1.4$~Gyr; Perlmutter et al. 1999), and from the shape of acoustic peaks in the cosmic microwave background spectrum ($14.0 \\pm 0.5$~Gyr; Knox, Christensen, \\& Skordis 2001).  One advantage of the thorium chronology is that, once the initial and current values of Th/$r$ in the star are specified, the age of the star depends only on the half-life of $^{232}$Th determined in the laboratory.  That is, one is not forced to invoke complex (and no doubt incomplete) models of Galactic chemical evolution, although insight from making simple attempts can still be obtained (e.g., Cowan et al. 1997). Here, $r$ denotes one of the stable $r$-process elements, often taken to be Eu. In the most simple model, the initial value of Th/$r$ may be regarded as that arising from a single supernova that exploded in the near vicinity (and perhaps may have even triggered the formation of) the presently observed star, e.g., CS~22892-052, early in the history of the Galaxy.  In fact, the large dispersion of Eu/Fe observed in halo stars (more than 2 orders of magnitude)  has been naturally explained by chemical evolution models that make such assumptions \\citep{Ishi99,Tsuj00}.  Thus far, the initial production of Th/$r$ has been determined by fitting theoretical nucleosynthesis results to the solar $r$-process pattern, with the assumption that the $r$-pattern was {\\it universal} in all astrophysical environments \\citep{Cowa99,Gori99}. However, a recent abundance analysis of CS~22892-052 with Keck~I (HIRES) has shown that the lighter neutron-capture elements ($Z \\le 48$) are distinctly underabundant with respect to the solar pattern, which is well fit by species with $Z \\ge 56$ \\citep{Sned00}. This implies that there likely exists at least two $r$-process sites that must be invoked to account for the entire set of elements associated with the $r$-process \\citep{Ishi00,Qian00}.  Furthermore, it should be noted that although Johnson \\& Bolte (2001) have derived an estimate of the minimum age of the universe, $11.0 \\pm 4.2$ Gyr, based on the abundances of thorium and other stable neutron-capture elements in five metal-poor stars, the estimated ages of these stars are very widely spread, from $\\approx 6$ to 18~Gyr (with overlapping error bars).  In addition, recent spectroscopic studies of several extremely metal-poor halo stars with SUBARU (HDS) has revealed that the spread of estimated chronometric ages from the Th-Eu pair appears to exceed the abundance uncertainties (Honda et al.  2002, in preparation).  Taken at face value, these results imply that the universality of the $r$-pattern may not hold over the actinides either.  The second discovered $r$-process-enhanced, extremely metal-poor star, CS~31082-001 (${\\rm [Fe/H]} = -2.9$) has provided a new, potentially quite powerful cosmochronometer, uranium \\citep{Cayr01,Hill02}. In principle, uranium might be expected to be a more precise chronometer than thorium owing to its shorter half life.  Cayrel et al. (2001) concluded that the age of this star is $12.5 \\pm 3$~Gyr, based on comparisons of the observed ratios of U/Th, U/Os, and U/Ir with their estimated production ratios \\citep{Cowa99,Gori99}. The discovery of uranium in CS~31082-001 has also prompted additional laboratory work to refine the oscillator strengths, for both uranium and thorium (Nilsson et al. 2002a, b), that are required to turn spectroscopic observations into measured abundances.  However, analysis of this star has also provided conclusive evidence that the $r$-process pattern is {\\it not universal} over the actinides; the ratio Th/Eu in this star is {\\it higher} than that of the solar $r$-process ratio.  CS~31082-001 would be younger than the solar system if the initial Th/Eu were taken to be universal, in conflict with the ages derived from ratios involving U and other stable $r$-process elements close to the third $r$-process peak.  Therefore, any age estimates that demand assumption of the universality of the $r$-process pattern may in fact be unreliable.  In addition to the above non-universality problem, the initial $r$-process pattern has thus far been determined theoretically by the superposition of nucleosynthesis results, where one is forced to assume constant temperatures, neutron number densities, and exposure times (or the number of captured neutrons) -- the most recent exploration of this approach is presented in Schatz et al. (2002).  These approximations have been necessary because of the lack of a reliable astrophysical model for the $r$-process site (at least that associated with the heavy $r$-process elements). At the moment, the neutrino-wind scenario is considered to be the most promising, ever since the work by Woosley et al. (1994), although alternative models involving neutron-star mergers have also shown some promise \\citep{Frei99b}. However, even the neutrino-wind model encounters serious problems, e.g., the high entropy ($\\sim 400 N_A k$) that led to the robust $r$-processing in the Woosley et al.  scenario has not been duplicated by other work \\citep{Taka94,Qian96}.  The purpose of this paper is, therefore, to construct a realistic $r$-process model based on the neutrino-wind scenario, in order to derive initial production ratios that might be useful for estimation of the minimum age of the universe, in particular by use of the U-Th chronometer pair.  Recently, Wanajo et al.  (2001) have demonstrated that the solar $r$-process pattern for $A \\approx 120-200$ is well reproduced in such a model from a compact proto-neutron star, owing to the inclusion of general relativistic effects. Although recent hydrodynamic studies have encountered difficulties attaining the required compactness, we use their model along with some improvements to be discussed in \\S~2.  Although some severe problems yet remain, it is clearly important to determine the initial $r$-process pattern based on site-specific, rather than only site-independent, models. In \\S~3 the $r$-process nucleosynthesis calculation is described, and the results are compared to the solar $r$-process pattern.  The initial electron fraction, $Y_e$, is varied from 0.39 to 0.49, but taken constant with time during the neutron-capture process.  In \\S~4, the age of CS~31082-001 is derived by comparison of the nucleosynthesis results obtained for each $Y_e$ with the abundance pattern of heavy elements in this star.  The implications of this study are discussed in \\S 5. ",
        "conclusions": "We have constructed a robust model of the $r$-process based on the neutrino-wind scenario for application to the U-Th cosmochronology. The model of neutrino winds was taken from Otsuki et al. (2000) and Wanajo et al. (2001), along with some improvements described above. The $r$-process with various neutron star masses ($M_{\\rm NS}$) and radii ($\\approx R_\\nu$) has already been explored extensively by Wanajo et al.  (2001). Thus, $M_{\\rm NS}$ and (asymptotic) $R_\\nu$ were fixed to be $2.0 M_\\odot$ and 10~km, respectively, for which the most robust $r$-processing was realized. On the other hand, the initial electron fraction $Y_e$ (assumed constant with time) was taken to be a free parameter, rather than fixed to be 0.40 as in Wanajo et al. (2001).  The model was tuned to fit the solar $r$-process pattern, since recent spectroscopic studies of extremely metal-poor stars have indicated a {\\it universal} (solar $r$-like) abundance pattern, at least between the second and third r-process peaks.  Attention was paid to reproduce the pattern near $A \\approx 195$, being close to Th and U in mass number. We introduced the freezeout temperature, $T_f$, as the most crucial parameter for constructing the model.  Comparing the nucleosynthesis results with the solar $r$-process pattern, $T_f$ was determined to be $1.0 \\times 10^9$~K. The nucleosynthesis yields were mass-integrated assuming simple evolutions of $L_\\nu$ and $R_\\nu$. The results between $A \\approx 120$ and 200 for $Y_e = 0.39-0.43$ models were in good agreement with the solar $r$-pattern, although the production of Th and U differed from model to model.  The results of our nucleosynthesis predictions were compared to the spectroscopic abundance pattern of the $r$-process-enhanced metal-poor star CS~31082-001.  Explorations of various combinations between the actinides and stable isotopes indicated that the $Y_e = 0.40$ case was found to be the best models for the age determination in this study. As a result, the age of CS~31082-001 was determined to be $14.1 \\pm 2.5$~Gyr by use of the U-Th chronometer, which can be regarded as a hard lower limit on the age of the universe. In fact, this age is in excellent agreement with the age of the universe derived by other dating techniques (e.g., globular clusters, Type Ia supernovae, and the cosmic-microwave background), as well as with other stellar radioactive age estimates using the Th-Eu \\citep{Sned96, Cowa97} and U-Th \\citep{Hill02, Scha02, Cowa02} pairs.  Such a strong constraint on the age of this star was only made possible by the use of U and Th in conjunction with one another.  Without uranium, the chronometers involving Th alone presented only rather loose constraints.  In this study the Pb abundance in CS~31082-001 was not available for age determination, since only an upper limit on its abundance has been presented to date. It should be noted that the predicted Pb abundance for $Y_e =0.40$ case is {\\it higher} than this upper limit, which seems difficult to solve \\citep{Scha02}. It is noteworthy, however, that freezeout effects can increase the ratio Th-U/Pb-Bi (see \\S~3), which should be examined further, as well as with other nuclear data sets. It remains possible that the solar $r$-process Pb abundance, inferred from the residual of the solar $s$-process components, is not correct \\citep{Arla99}. Clearly, re-investigation of the Pb abundance in CS~31082-001, as well expansion of the searches for additional $r$-process-enhanced metal-poor stars, in which both U and Th are measurable, are highly desirable in the near future.  It is also of importance to obtain measurements of the Pt abundance in CS~31082-001 (which is only possible with the strong Pt absorption features observable in the near UV from space), so as to constrain the height, as well as the position, of the third peak \\citep{Cowa02}.  We emphasize that the conclusions in this study should be considered as merely the first step toward a better understanding of the dating technique involving the U-Th chronometer. In particular, we note that the inferred age of CS~31082-001 could be significantly changed if we had utilized a different nuclear mass model, as demonstrated by Griely and Arnould (2001). We expect, however, that a detailed comparison of the nucleosynthesis results with the solar $r$-process pattern, as in \\S~3, can distinguish better mass models. We suggest that the freezeout temperature, $T_f$ (or the asymptotic temperature in more realistic models), is a useful parameter, as the freezeout effects are of importance to the final $r$-process pattern. While the waiting-point approximation, often referred to as the ``classical $r$-process model,'' has been an effective (and thus far, the only) method used to test the nuclear mass models, we should keep in mind that non-equilibrium processes can lead to non-negligible effects, as discussed in \\S~3 (e.g., the positions of the $r$-process peaks, the smoothness of abundance pattern, and the ratio Th-U/Pb-Bi).  It should also be noted that we would not insist that the neutrino-wind model explored in this study is necessarily the best (or only) astrophysical site for the $r$-process. In fact, the physical conditions in this model, namely, the rather massive proto-neutron star ($2.0 M_\\odot$) with rapid contraction of $R_\\nu$ to be 10~km, seem difficult to be achieved according to the current hydrodynamic studies. Clearly, more study is needed to seek the true $r$-process site. Nevertheless, it is encouraging that the age of the CS~31082-001 determined in this study is in excellent agreement with the age of the universe derived by other dating techniques. We hope our results provide some fresh insights to the future modeling of the true astrophysical $r$-process site."
    },
    "0206/hep-ph0206010_arXiv.txt": {
        "abstract": "High Energy Neutrino signals from PeVs up to  GZK, $E_{\\nu}\\geq 10^{19}$ eV, energies  amplified as Upward and Horizontal Tau air Showers,\\textbf{UPTAUS}, and \\textbf{HORTAUS}, may flash toward mountain peaks, airplanes, balloons  and satellites in future Tau Air-Showering experiments. Observing from  high mountains the detectable crown masses at UHE $\\nu$ energies around EeVs  are comparable to a few $km^3$, while observing from higher balloons or satellites at orbit altitudes, at GZK  energies  , the Horizontal Crown effective Masses  may even exceed $150$ $km^3$. The consequent event rate by expected GZK or Z-Showering neutrinos models must arise at large rate by present scheduled (EUSO-like) or future (Crown Array Detector) satellite in Space. ",
        "introduction": "Rarest TeVs gamma signals are at present the most extreme and rarest trace of  High Energy Astrophysics. The TeVs signals have shown new power-full Jets blazing to us from Galactic or extragalactic edges. At PeVs energies astrophysical Gamma cosmic rays should also be present, but, excluded a very rare and elusive Cyg$X3$ event, they are  not longer being observed. While the corresponding PeVs charged cosmic rays are abundantly hitting the atmosphere, these missing PeVs gamma sources are very probably mostly absorbed by their photon interactions (photo-pion productions, electron pairs creation) at the source environment and/or along the photon propagation into the cosmic Black Body Radiation (BBR) or into other diffused (radio,infrared,optical) background radiation. Unfortunately PeVs charged cosmic rays, bend and bounded in a  random walk by Galactic magnetic fields, loose their original directionality and their astronomical relevance; their resident time in the galaxy is much longer ($\\geq 10^{3}$-$10^{5}$) than neutral ones, as gamma rays, making the charged cosmic rays more probable to be observed by nearly a comparable ratio. On the contrary astrophysical UHE neutrino signals at $10^{13}$eV-$10^{19}$eV (or higher GZK energies) are unaffected by any radiation cosmic opacity and may open a very new exciting window to High Energy Astrophysics. Lower energy astrophysical UHE $\\nu$ at $10^{9}$eV-$10^{12}$eV should also be present, but their signals are (probably) drowned  by the dominant diffused atmospheric $\\nu$ secondaries noises produced by the same charged (and smeared) UHE cosmic rays (while hitting terrestrial atmosphere), the so called atmospheric neutrinos. In a very far corner, at lowest (MeVs) energy  windows, the abundant and steady solar neutrino flux  and the prompt (but rarer) neutrino burst from a nearby Super-Novae (SN 1987A), have been in last twenty years, already successfully explored. \\subsection{The UHE $\\nu$ $Km^3$ detectors} The UHE $10^{13}$eV-$10^{16}$eV $\\nu$ 's  astronomy, being weakly interacting and rarer, may  be captured mainly inside huge volumes, bigger than Super-Kamiokande ones; at present  most popular detectors consider underground ones (Cubic Kilometer Size like AMANDA-NESTOR)\\cite{Halzen2002} or (at higher energy $10^{19}$eV-$10^{21}$eV) the widest Terrestrial atmospheric sheet volumes (Auger-Array Telescope or EUSO atmospheric Detectors). Underground $km^3$ detection is based mainly on $\\nu_{\\mu}$ (above hundred TeVs energies, after their interaction with matter) leading to $\\mu$ kilometer size lepton tracks \\cite{Gandhi et al 1998}. Rarest atmospheric horizontal shower are also expected by $\\nu$ interactions in air (and, as we shall discuss, in the Earth Crust). While $km^3$ detectors are optimal for PeVs neutrino muons, the Atmospheric Detectors (AUGER-EUSO like) exhibit a minimal threshold at highest ($\\geq 10^{18} eV$) energies. \\subsection{ UHECR EeV $n$ and PeV$\\nu$  Astronomy} Incidentally just around such EeV ($10^{18} eV$) energies an associated Ultra High Energy Neutron Astronomy  might be possible because of the relativistic neutrons boosted lifetime, comparable to our distance from the Galactic center. Therefore UHE  neutrons at EeV may be source of part of UHECR data. Indeed a $4\\%$ galactic anisotropy and clustering in EeV cosmic rays is recently emerging \\cite{AGASA 1999} along our nearby galactic plane. Therefore AGASA might have already experienced a first UHECR-Neutron astronomy \\textbf{UHENA}. This EeV-UHENA signals must also be source of a parasite ($10^{17}-10^{16}$ eV) secondary tails of UHE neutrino $\\bar{\\nu_{e}}$  from the same neutron beta decay in flight. Their oscillations in galactic or extragalactic flights (analogous to atmospheric and solar ones) must guarantee the presence of all lepton flavours nearly at equal foot: $\\bar{\\nu_{e}}$,$\\bar{\\nu_{\\mu}}$ $\\bar{\\nu_{\\tau}}$. The latter UHE $\\bar{\\nu_{\\tau}}$ imprint (added to other local astrophysical UHE $\\nu$ production) could be already recorded \\cite{Fargion2000-2002} as Upward Tau Air-Showers  Terrestrial Gamma Flash (assumed secondaries $\\gamma$ of Upward Tau Air-Showers and Horizontal Tau Air-Showers). These estimates are part of present article below. \\subsection{ GZK $\\nu$ Astronomy and Z-Showering} Finally,  let us also remind that at highest energy edges ($\\geq 10^{19}-10^{20} eV$), a somehow correlated New UHE Astronomy is also expected for charged Cosmic Rays; indeed these UHECR have such a large rigidity to avoid any bending by random galactic or extragalactic magnetic fields;  being nearly undeflected UHECR should point toward the original sources showing in sky a new astronomical map. (There have been very reiterated attempts to consider extreme galactic or extragalactic coherent magnetic fields able to bend the well known nearby, like Virgo Clusters, M $87$ AGN, sources in a diffused map, however their needed magnetic energy density and coherent lenghts are in our opinion totally unrealistic). Moreover such UHECR astronomy is bounded by the ubiquitous cosmic $2.75 K^{o}$ BBR screening (the well known Greisen,Zat'sepin,Kuzmin GZK cut-off) limiting its origination inside a very local ($\\leq 20 Mpc$) cosmic volume. Surprisingly, these UHECR above GZK (already up to day above $60$ events) are not pointing toward any known nearby candidate source. Moreover their nearly isotropic arrival distributions underlines and testify a very possible cosmic origination, in disagreement with any local (Galactic plane or Halo) expected footprint by GZK cut-off. A very weak Super-Galactic imprint seem to be present but already above GZK volume.  This opened a very hot debate in modern astrophysics known as the GZK paradox. Possible solutions has been found recently beyond Standard Model assuming a non-vanishing neutrino mass. Indeed at such Ultra-High energies, neutrino at ZeV energies ($\\geq 10^{21}eVs$) hitting onto  relic cosmological light ($0.1-4 eV$ masses) neutrinos \\cite{Dolgov2002} nearly at rest in Dark Hot Halos (galactic or in Local Group) has the unique possibility to produce UHE resonant Z bosons (the so called Z-burst or better Z-Showering scenario): \\cite{Fargion Salis 1997}, \\cite{Fargion Mele  Salis 1999}, \\cite{Yoshida  et all 1998},\\cite{Weiler 1999}; for a more updated scenario see \\cite{Fargion et all. 2001b}. Indeed UHE neutrinos are un-effected by magnetic fields and by BBR screening; they may reach us from far cosmic edges with negligible absorption. The UHE Z-shower in its ultra-high energy nucleonic secondary component may be just the observed final UHECR event on Earth. This possibility has been reinforced by very recent correlations (doublets and triplets events) between UHECR directions with brightest Blazars sources at cosmic distances (redshift $\\geq 0.1$) quite beyond ($\\geq 300 Mpc$) any allowed GZK cut-off \\cite{Tinyakov-Tkachev2001} \\cite{AGASA 1999},\\cite{Takeda et all}. Therefore there might be a role for GZK neutrino fluxes, either as primary in the Z-Showering scenario or, at least, as (but at lower intensities) necessary secondaries of all those UHECR primary absorbed in cosmic BBR radiation fields. Naturally other solutions as topological defects or primordial relics decay may play a role as a source of UHECR, but the observed clustering \\cite{AGASA 1999}, \\cite{Tinyakov-Tkachev2001},\\cite{Takeda et all}  seems to favor compact sources. The most recent evidence for self-correlations clustering at $10^{19}$, $2\\cdot10^{19}$, $4\\cdot10^{19}$ eVs energies observed by AGASA (Teshima, ICRR26 Hamburg presentation 2001) maybe a first reflection of UHECR Z-Showering secondaries: $p, \\bar{p}, n, \\bar{n}$ \\cite{Fargion et all. 2001b}. A very recent solution beyond the Standard Model (but within Super-Symmetry) consider Ultra High Energy Gluinos as the neutral particle bearing UHE signals interacting nearly as an hadron in Terrestrial Atmosphere \\cite{Berezinsky 2002}; this solution has a narrow window  for gluino masses allowable (and serious problems in production bounds), but it is an alternative that deserves attention. To conclude the puzzle one finally needs to scrutiny the UHE $\\nu$ astronomy and to test the GZK solution within Z-Showering Models by any independent search on Earth for such UHE neutrinos traces above PeVs reaching either EeVs-ZeVs extreme energies. \\subsection{UHE $\\nu$ Astronomy by $\\tau$ Air-Shower} Recently \\cite{Fargion et all 1999},\\cite{Fargion2000-2002} it has been proposed a new competitive (and in our opinion, more convenient) UHE $\\nu$ detection based on ultra high energy $\\nu_{\\tau}$ interaction in matter and its consequent secondary $\\tau$ decay in flight while escaping from the rock (Mountain Chains) or water (Sea)  in air leading to Upward or Horizontal $\\tau$ Air-Showers (UPTAUs and HORTAUs),\\cite{Fargion2001a},\\cite{Fargion2001b}.  In a pictorial way one may compare the UPTAUs and HORTAUs as the double bang processes expected in $km^3$ ice-water volumes \\cite{Learned Pakvasa 1995} : the double bang is due first to the UHE $\\nu_{\\tau}$ interaction in matter and secondly by its consequent $\\tau$ decay in flight. Here we consider  a (hidden) UHE $\\nu$-N Bang $in$ (the rock-water within a mountain or the Earth Crust) and a $\\tau$ bang $out$ in air, whose shower is better observable at high altitudes. The main power of the UPTAUs and HORTAUs detection is the huge amplification of the UHE neutrino signal, which may deliver almost all its energy in numerous secondaries traces (Cherenkov lights, gamma, X photons, electron pairs, collimated muon bundles). Indeed the multiplicity in $\\tau$ Air-showers secondary particles, $N_{opt} \\simeq 10^{12} (E_{\\tau} / PeV)$, $ N_{\\gamma} (< E_{\\gamma} > \\sim  10 \\, MeV ) \\simeq 10^8 (E_{\\tau} / PeV) $ , $N_{e^- e^+} \\simeq 2 \\cdot 10^7 ( E_{\\tau}/PeV) $ , $N_{\\mu} \\simeq 3 \\cdot 10^5 (E_{\\tau}/PeV)^{0.85}$ makes easy the UPTAUs-HORTAUs discover. These HORTAUs, also named Skimming neutrinos \\cite{Feng et al 2002}, studied also in peculiar approximation in the frame of AUGER experiment, \\cite{Bertou et all 2002}, maybe also originated on front of mountain chains \\cite{Fargion et all 1999}, \\cite{Hou Huang 2002}  either by $\\nu_{\\tau}N$, $ \\bar\\nu_{\\tau}N$ interactions as well as by $ \\bar\\nu_{e} e \\rightarrow W^{-} \\rightarrow \\bar\\nu_{\\tau} \\tau$. This new UHE $\\nu_{\\tau}$ detection is mainly based on the oscillated UHE neutrino $\\nu_{\\tau}$ originated by more common astrophysical $\\nu_{\\mu}$, secondaries of pion-muon decay at PeVs-EeVs-GZK energies. These oscillations are guaranteed  by Super Kamiokande evidences for flavour mixing within GeVs atmospheric neutrino data as well as by most solid and recent evidences of complete solar neutrino mixing observed by SNO detector. HORTAUs from mountain chains must nevertheless occur, even for no flavour mixing, as being inevitable $\\bar\\nu_{e}$ secondaries of common pion-muon decay chains ($\\pi^{-}\\rightarrow \\mu^{-}+\\bar\\nu_{\\mu}\\rightarrow e^{-}+\\bar\\nu_{e}$) near the astrophysical sources at Pevs energies. They are mostly absorbed by the Earth and are only rarely arising as UPTAUS. Their Glashow resonant interaction allow them to be observed as HORTAUs only within a very narrow and nearby crown edges at horizons (not to be discussed here). At wider energies windows ($10^{14}eV- 10^{20}eV$) only neutrino $\\nu_{\\tau}$, $\\bar{\\nu}_{\\tau}$ play a key role in UPTAUS and HORTAUS. These Showers might be easily detectable looking downward the Earth's surface from mountains, planes, balloons or satellites observer. Here the Earth itself acts as a \"big mountain\" or a wide beam dump target. The present upward $\\tau$ at horizons should not be confused with an independent and well known, complementary (but rarer) Horizontal Tau Air-shower originated inside the same terrestrial atmosphere: we shall referee to it as the Atmospheric Horizontal Tau Air-Shower. The same UPTAUS have a less competitive upward showering due to $\\nu_{e}$ $\\bar\\nu_{e}$ interactions with atmosphere, showering in thin upward air layers \\cite{Berezinsky 1990}: let us label this atmospheric Upward Tau as A-UPTAUs  and consider its presence  as a very small additional contribute, because rock is more than $3000$ times denser than air. Therefore at different heights we need to estimate the UPTAUS and HORTAUs event rate occurring along the thin terrestrial crust below the observer, keeping care of their correlated  variables: from a very complex sequence of functions we shall be able to define and evaluate the effective HORTAUs volumes keeping care of the thin shower beaming angle, atmosphere opacity and detector thresholds. At the end of the study, assuming any given neutrino flux, one might be easily able to estimate at each height $h_{1}$ the expected event rate and the ideal detector size and sensibility for most detection techniques (Cherenkov, photo-luminescent, gamma rays, X-ray, muon bundles). ",
        "conclusions": "The discover of the expected UHE neutrino Astronomy is urgent and just behind the corner. Huge volumes are necessary. Beyond underground $km^3$ detectors a new generation of UHE neutrino calorimeter lay on front of mountain chains and just underneath our feet: The Earth itself  offers huge Crown Volumes as Beam Dump calorimeters observable via upward Tau Air Showers, UPTAUs and HORTAUs. Their effective Volumes as a function of the quota $h_1$ has been derived by an analytical function variables in equation $(17), (20),(23),(28)$ and Appendix B. These Volumes are discussed below and summirized in the last column of next Table (Fig.9) at the end of the paper. At a few tens meter altitude the UPTAUs and HORTAUs Ring are almost overlapping. At low altitude $h_{1}\\leq 2 km$ the HORTAUs are nearly independent on the $\\rho$ matter density: ${\\Delta M_{eff.}(h_{1}=2 km)(\\rho_{Water})}= 0.987 km^3 $ ${\\Delta M_{eff.}(h_{1}=2km)(\\rho_{Rock})}= 0.89 km^3 $ These volumes are the effective Masses expressed in Water equivalent volumes. On the contrary at higher quotas, like highest Mountain observations sites, Airplanes, Balloons and Satellites, the matter density of the HORTAUs Ring Areas play a more and more dominant role asymptotically $proportional$ to ${\\rho}^{-2}$: ${\\Delta M_{eff.}(h_{1}=5 km)(\\rho_{Water})}= 3.64 km^3 $ ${\\Delta M_{eff.}(h_{1}=5 km)(\\rho_{Rock})}= 2.17 km^3 $ From Air-planes or balloons the effective volumes $M_{eff.}$ increases and the density $\\rho$ plays a relevant role. ${\\Delta M_{eff.}(h_{1}=25 km)(\\rho_{Water})}= 20.3 km^3 $ ${\\Delta M_{eff.}(h_{1}=25 km)(\\rho_{Rock})}= 6.3 km^3 $ Finally from satellite altitudes the same effective volumes $M_{eff.}$ are reaching extreme values. \\begin{displaymath} {\\Delta M_{eff.}(h_{1}=500 km)(\\rho_{Water})}= 150.6 km^3 \\end{displaymath} \\begin{displaymath} {\\Delta M_{eff.}(h_{1}=500 km)(\\rho_{Rock})}=30 km^3 \\end{displaymath} These Masses must be compared with other proposed $km^3$ detectors, keeping in mind that these HORTAUs signals conserve the original UHE $\\nu$ direction information within a degree. One has  to discriminate  HORTAUS (only while observing from satellites) from Horizontal High Altitude Showers (HIAS) \\cite{Fargion2001b}, due to rare UHECR showering on high atmosphere. One might also remind the UPTAUs (at PeVs energies) volumes as derived in Appendix B and in \\cite{Fargion2000-2002} whose values (assuming an arrival angle$\\simeq 45^o- 60^o$ below the horizons) are nearly $proportional$ to the $\\rho$ density: \\begin{displaymath} {\\Delta M_{eff.}(h_{1}=500 km)(\\rho_{Water})}= 5.9 km^3 \\end{displaymath} \\begin{displaymath} {\\Delta M_{eff.} (h_{1}=500 km)(\\rho_{Rock})}=15.6 km^3 \\end{displaymath} These widest Masses values, here estimated analytically for any quota, are offering an optimal opportunity to reveal UHE $\\nu$ at PeVs and EeVs-GZK energies by crown array detectors (scintillators, Cherenkov, photo-luminescent) facing vertically the Horizontal edges, located at high mountain peaks or at air-plane low sides and finally on  balloons and satellites. As it can be seen in last column of final Table , the ratio $ R $ between HORTAUs events and Showers over atmospheric UHE $\\nu$ interaction is a greater and greater number with growing height, implying a dominant role (above two order of magnitude) of HORTAUS grown in Earth Skin Crown over Atmospheric HORTAUs."
    },
    "0206/astro-ph0206180_arXiv.txt": {
        "abstract": "{We have computed accurate 1-D solar models including both a macroscopic mixing process in the solar tachocline  as well as up-to-date microscopic physical ingredients. Using sound speed and density profiles inferred through primary inversion of the solar oscillation frequencies coupled with the equation of thermal equilibrium, we have extracted the temperature and hydrogen abundance profiles.  These inferred quantities place strong constraints on our theoretical models in terms of the extent and strength of our macroscopic mixing, on the photospheric heavy elements abundance, on the nuclear reaction rates such as $S_{11}$ and $S_{34}$ and on the efficiency of the microscopic diffusion. We find a good overall agreement between the seismic Sun and our models if we introduce a macroscopic mixing in the tachocline and allow for variation within their uncertainties of the main physical ingredients. From our study we deduce that the solar hydrogen abundance at the solar age is $X_{\\rm inv}=0.732\\pm 0.001$ and that based on the $^9$Be photospheric depletion, the maximum extent of mixing in the tachocline is 5\\% of the solar radius. The nuclear reaction rate for the fundamental $pp$ reaction is found to be $S_{11}(0)=4.06\\pm 0.07$ $10^{-25}$ MeV barns, i.e., 1.5\\% higher than the present  theoretical determination. The predicted solar neutrino fluxes are discussed in the  light of the new SNO/SuperKamiokande results.} ",
        "introduction": "Over the past decade our understanding of the solar interior has improved significantly. Today with the precise helioseismic data available from the GONG (Global Oscillation Network Group) ground based instruments and the SOHO (SOlar and Heliospheric Observatory) space experiments (Gough et al.~1996; Thompson et al.~1996; Fr\\\"ohlich et al. 1997; Gabriel et al. 1997; Schou et al.~1998), the detailed internal structure and complex dynamics of our star can be inferred with reasonable accuracy using inversion techniques.  In addition to sound speed $c$ and density $\\rho$ profiles, the internal rotation rate $\\Omega$ can also be inferred. It reveals that on the top of an almost uniformly rotating radiation zone (with a rotation period of about 28 days), the bulk of the convection zone is differentially rotating with properties close to what is deduced from sunspot tracking, i.e., a period at the equator of 25 days and at the pole of 33 days corresponding to a contrast $\\Delta\\Omega$ of 30\\%. The sharp transition region between these two distinct zones, located around $0.7R_\\odot$, has been called the tachocline (Spiegel \\& Zahn 1992); it is thought to play an important role in determining the structure and the chemical evolution of the Sun (Brun, Turck-Chi\\`eze \\& Zahn 1999; Elliott \\& Gough 1999).  Indeed, it appears that the solar structure deduced from helioseismology and the observed photospheric compositions can not be explained adequately without invoking some mixing in the radiative interior (Brun et al.~1999). This conclusion has been drawn after a careful study of the microscopic processes present in solar models (\\jcd\\ et al.~1996; Morel, Provost \\& Berthomieu 1997; Bahcall, Basu \\& Pinsonneault 1998a; Brun, Turck-Chi\\`eze \\& Morel 1998).  More precisely, in the early 90's, after significant improvements in the description of the solar plasma through better equation of state, opacities and nuclear reaction rates, helioseismic studies  have established the need for microscopic diffusion of helium and heavy elements in the radiative interior (\\jcd\\ et al.~1993).  But it was soon realized that models including only microscopic diffusion exhibit sharp composition gradients below the base of the convection zone which are not consistent with helioseismic data; these favour instead smoother composition profiles within this region (Basu \\& Antia 1994). In spite of further improvements in solar models, this discrepancy still persists around  0.7 $R_{\\odot}$  (\\jcd\\ et al.~1996; Brun et al.~1998), suggesting that some extra mixing must be implemented in the models.  Further, using primary inversions for sound speed and density and the equations of thermal equilibrium, Antia \\& Chitre~(1998) have inferred the hydrogen abundance profile in the radiative interior. This profile confirms the presence of such mixing in the Sun, as the hydrogen abundance appears to be almost constant in the region $r>0.68R_\\odot$.  Another evidence for mixing occurring in that region comes from the photospheric light elements composition. Purely microscopic processes cannot reproduce the under abundance of lithium observed in the Sun and in open clusters (Grevesse, Noels \\& Sauval 1996; Cayrel 1998; Richard et al.~1996; Turcotte et al.~1998; Brun et al.~1999).  We are thus compelled to introduce some mixing processes in the stably stratified radiative interior.  The possible causes of instabilities leading to such mixing are the solar rotation, the magnetic field or penetrative convection (Zahn 1998).  The recent study of Balachandran \\& Bell (1998) on the photospheric light elements abundance of $^7$Li and $^9$Be puts strong constraints on the extent, amplitude and location of such instabilities. It is now believed that only the lithium is significantly depleted,  by more than a factor of 100 in comparison to the meteoritic composition, while the beryllium has varied by only 10\\% over the last 4.6 Gyr. The temperatures at which these two species are destroyed by nuclear burning are respectively $\\sim 2.7\\times10^6$ K (at $0.66 R_{\\odot}$) and $\\sim 3.2\\times10^6$ K (at $0.59 R_{\\odot}$), which are relatively close to the temperature at the base of the convection zone $\\sim 2.2\\times10^6$ K at $0.713 R_{\\odot}$ (\\jcd\\ et al.~1991). This implies that any macroscopic processes for such lithium destruction have to be located near the top of the radiation zone and cannot extend deeper than ($\\sim 8\\%$) in solar radius without producing an excessive destruction of $^9$Be.  This requirement is satisfied if the mixing is confined in the tachocline. For this reason, Brun et al.~(1999) calculated the mixing occurring in that layer,  based on Spiegel \\& Zahn's hydrodynamical description of the tachocline.  They found indeed that such mixing improved the agreement between the models and the Sun, provided the secular variation of the tachocline was taken into account. Alternative approaches based on gravity waves (Montalban \\& Schatzman 1996) and magnetic field (Barnes, Charbonneau \\& MacGregor 1999) have also been studied.  In this paper we intend to go further in understanding the influence  of tachocline mixing on the solar internal structure and composition by using the powerful seismic diagnostic. In particular, we examine the cumulative effect of varying, within their intrinsic uncertainties, the basic microscopic ingredients such as the nuclear reaction cross sections, the microscopic diffusion coefficients, etc. Seismic probing of the solar structure has been improved to a level where it can be used to constrain physical processes, since the Sun provides a ready-made cosmic laboratory for testing various aspects of physics. For example, there have been some attempts to constrain the nuclear reaction rate for {\\it pp} reaction using helioseismic data (Antia \\& Chitre 1998, 1999; Degl'Innocenti, Fiorentini \\& Ricci~1998; Schlattl, Bonanno \\& Paterno~1999). These studies indicate that the cross-section for {\\it pp} reaction needs to be increased by a few percent over the currently accepted value (Adelberger et al.~1998). Using a similar approach Weiss, Flaskamp \\& Tsytovich (2001) found that enhancing the electron screening by about 5\\% improves the agreement between  solar model and helioseismically inferred sound speed. Therefore, we would like to revise, among other quantities, the $pp$ cross-section deduced from helioseismology, by using in our solar models either the weak (Salpeter 1954) or intermediate (Mitler 1977) treatment for electron screening. We also deduce the photospheric  hydrogen abundance $X_{\\rm ph}$, the maximum extent $h$ of the tachocline mixing allowed in a solar model and predict the theoretical neutrino fluxes in light of the recent SNO results (Ahmad et al.~2001, 2002).  We have organised our paper as follows. In section 2, we briefly recall how we compute our 1-D solar models with or without the presence of tachocline mixing, describe our inversion techniques of the solar acoustic frequencies and demonstrate the need for further progress in solar modelling.  In section 3, we present our latest results on the sound speed, density, hydrogen abundance and temperature profiles obtained with our modified solar tachocline models and discuss the resulting neutrino fluxes. Finally in section 4, we comment on our findings and outline our conclusions. ",
        "conclusions": "In this work our goal was to assess the effect of uncertainties in nuclear reaction rates, atomic data and diffusion coefficients on solar models, which we have compared with the results of helioseismic inversions. We have concluded in \\S \\ref{ism} that a variation of the nuclear reaction rates can have a significant impact on the solar structure and that current modelling coupled with seismic data favours some change of the accepted central value of three important nuclear reaction rates, {\\it pp}, $^3$He-$^3$He and $^3$He-$^4$He. While the {\\it pp} reaction rates needs to be increased by about 1.5\\% over the currently accepted value from Adelberger et al.~(1998), the constraints on other reaction rates are less clear. Even the increase of {\\it pp} rate involved in this study is less than that inferred by Antia \\& Chitre~(1999) and Degl'Innocenti et al.~(1998) and the difference can be attributed to different treatment of nuclear energy generation in CESAM as compared to the version of Bahcall's energy routine used in earlier estimates. More precisely, most of the discrepancy comes from the different screening prescriptions used. In the previous study of Antia \\& Chitre~(1999), the screening formulation of Graboske et al.~(1973) was assumed as opposed to Salpeter (1954) or Mitler (1977) in this new study. We refer to Dzitko et al.~(1995) and Wilets et al. (2000) for a detailed comparison of several weak, intermediate and strong, screening prescriptions.  Our work confirms the result of Brun et al.~(1999) that the implementation of macroscopic mixing in the tachocline improves the agreement between solar models and seismic Sun (\\S 3.2). In particular, direct comparison of hydrogen abundance profiles between our models and the Sun as inferred from seismic inversions has demonstrated that the models with tachocline mixing are in much better agreement with the Sun in tachocline region. Comparison of hydrogen abundance profile in solar models with inferred profiles show that there is still a discrepancy of about 0.003 (\\S 3.3). The largest discrepancies occur in the region close to the tachocline and at about $r=0.2R_\\odot$, where the sound speed also shows maximum discrepancy. Thus it appears that there is still some scope for improving the formulation for calculating the mixing in tachocline region. We were not able to achieve any significant improvement compared to the results of  Brun et al.~(1999) by adjusting the parameters of the tachocline model. With some modifications the agreement improves in the tachocline region, like in model $NM$, but it tends to worsen in other places. We have found as well that based on the $^9$Be photospheric depletion, the maximum extent of the mixing in the tachocline is 5\\% of solar radius. It is quite possible that a major part of the remaining discrepancies (about 0.1\\% in sound speed, 1\\% in density and 0.003 in $X$) in our improved models may be due to uncertainties in input physics, like the opacities or equation of state.  Comparing the surface hydrogen abundance in solar models and those obtained by inversions (Table~3) with the abundance inferred directly from seismic data (Basu 1998), we again find a difference of 0.001. This suggests that discrepancies of this order are present in the input physics which is used in inferring these independent measurements of surface hydrogen abundance. The helium abundance drawn from the structure of the ionisation zones is particularly sensitive to the equation of state (Basu \\& Antia 1995), while its abundance inferred from the solar models depends on the whole input physics used in constructing those. Similarly, the seismically inferred hydrogen abundance profile is sensitive to opacity as well as to the $Z$ profile used in inversion. Thus the discrepancy between these three independent estimates of surface hydrogen abundance probably gives an estimate of errors in input physics. The difference between solar models with tachocline mixing and seismically inferred profiles is of the same order and could be due also to remaining uncertainties in input physics.  As seen in Fig.~1 the evolutionary solar models show a significant departure from seismic inferences inside the convection zone. We find that scaling the radius of these models by 1.0003 before taking the differences with inverted profile removes most of the discrepancy in the convection zone. We believe that this is due to uncertainties in the treatment of surface layers. In the solar models the surface is defined by the layer where the temperature equals effective temperature and it is quite possible that because of uncertainties in surface layers the location of this point has an error of about 200 km in our evolutionary models. We would like to point out that adjusting the solar radius in the model by 200 km does not remove this discrepancy, since even in the new model the position of the surface has the same uncertainty. Thus this scaling of radius has no relation to uncertainty in solar radius itself.  The calculated neutrino fluxes in solar models with tachocline mixing are found to be somewhat lower than those in standard solar model of Bahcall et al.~(2001), but within the $1\\sigma$ error limits (\\S 3.4). The main reason for reduction is the increase in {\\it pp} reaction rate, which reduces the central temperature required to generate the solar luminosity. Recently, using both the charged and neutral current channels, SNO has measured the total $^8$B neutrino flux of $5.09\\times10^6$ cm$^{-2}$ s$^{-1}$ (Ahmad et al.~2002). This is somewhat larger than the value we find in our models.  In the modified models $N1, N2$ the neutrino fluxes are much lower, because we have chosen to modify the nuclear reaction rates to reduce these fluxes. If instead as with model $N03$ we had modified $S_{33}$ and $S_{34}$ in opposite directions the neutrino fluxes would have increased and it would be possible to get values close to that inferred by Ahmad et al.~(2002). Also our increase of {\\it pp} reaction rate by 3.5\\% in some of these models is probably an overestimate, since seismic models with same input physics appear to need an increase by only 1.5\\% to produce the required solar luminosity. When almost all the sources of uncertainties are taken into account in evaluating the theoretical neutrino fluxes, our result is within 1$\\sigma$ of the other published theoretical neutrino fluxes and the recent SNO compilation. It will be interesting in the near future, when the SNO collaboration results will have been integrated over a longer time to see if the helioseismic tool will be able to constrain even more efficiently the main nuclear cross sections, via a careful analysis of the neutrino spectrum, and also delineate the parameter space for mass-squared difference-mixing angle plane.  \\vspace{-0.4cm} \\begin{acknowledgement} We thank the anonymous referee for constructive comments and suggestions which have led to a considerable improvement in the final version of the paper. ASB would like to thank S. Turck-Chi\\`eze for useful discussions and the Service d'Astrophysique, for access to their computers. This work was partly supported by NASA through grants NAG5-2256 and NAG5-8133. SMC is grateful to DAE-BRNS for support under the Senior Scientist Scheme. \\end{acknowledgement}"
    },
    "0206/astro-ph0206463_arXiv.txt": {
        "abstract": "{ We report the detection of H$_2$, C\\,{\\sc i}, C\\,{\\sc i}\\,$^\\star$, C\\,{\\sc i}\\,$^{\\star\\star }$ and Cl\\,{\\sc i} lines in a near Solar-metallicity ([Zn/H$]=-0.13$) damped Lyman-$\\alpha$ (DLA) system at $z_{\\rm abs}=1.962$ observed on the line of sight to the quasar Q\\,0551$-$366. The iron-peak elements, ${\\rm X}={\\rm Fe}$, Cr and Mn are depleted compared to zinc, [X/Zn$]\\sim -0.8$, probably because they are tied up onto dust grains. Among the three detected H$_2$-bearing clouds, spanning 55 km s$^{-1}$ in velocity space, we derive a total molecular hydrogen column density $N($H$_2)=2.6\\times 10^{17}$ cm$^{-2}$ and a mean molecular fraction $f=2N($H$_2)/(2N($H$_2)+N($H\\,{\\sc i}$))=1.7\\times 10^{-3}$. The depletion of heavy elements (S, Si, Mg, Mn, Cr, Fe, Ni and Ti) in the central component is similar to that observed in the diffuse neutral gas of the Galactic halo. This depletion is approximately the same in the six C\\,{\\sc i}-detected components independently of the presence or absence of H$_2$. The gas clouds in which H$_2$ is detected always have large densities, $n_{\\rm H}>30$ cm$^{-3}$, and low temperatures, $T_{\\rm 01}\\la 100$ K. This shows that presence of dust, high particle density and/or low temperature are required for molecules to be present. The photo-dissociation rate derived in the components where H$_2$ is detected suggests the existence of a local UV radiation field similar in strength to the one in the Galaxy. Star formation therefore probably occurs near these H$_2$-bearing clouds. ",
        "introduction": "Damped Lyman-$\\alpha$ (DLA) absorption line systems observed in the spectrum of quasars are associated with large H\\,{\\sc i} column densities ($N$(H\\,{\\sc i}$)>2\\times 10^{20}$ cm$^{-2}$). Therefore, they are probably associated with regions of the Universe where star formation occurs (Pettini et al. 1997). It is still unclear whether the gas producing the absorption lines is located in large, fast-rotating protogalactic discs (Wolfe et al. 1986, Prochaska \\& Wolfe 1997), in interacting building blobs (Haehnelt et al. 1998) or in density fluctuations of galaxy haloes (Ledoux et al. 1998).  Whatever the exact nature of DLA systems may be, molecules (especially H$_2$) are expected to be found in those clouds. However, despite intensive searches (e.g. Black et al. 1987, Ge \\& Bechtold 1999, Petitjean et al. 2000) only four detections of H$_2$ have been reported to date. Recently, a fifth case has been discovered serendipitously by Levshakov et al. (2002) at $z_{\\rm abs}=3.025$ toward Q\\,0347$-$383.  The most recent survey for molecular hydrogen in 11 $z_{\\rm abs}>1.8$ DLA systems, capitalizing on the unique capabilities of the UVES high-resolution spectrograph of the ESO Very Large Telescope, has given stringent upper limits, in the range $1.2\\times 10^{-7}$ -- $1.6\\times 10^{-5}$, for the molecular fraction $f=2N($H$_2)/(2N($H$_2)+N($H\\,{\\sc i}$))$ in nine of the systems (Petitjean et al. 2000). Two possible detections have also been reported at $z_{\\rm abs}=2.374$ toward Q\\,0841$+$129 (Petitjean et al. 2000) and at $z_{\\rm abs}=3.390$ toward Q\\,0000$-$263 (Levshakov et al. 2000) based, however, on the detection of only two weak features located into the Lyman-$\\alpha$ forest. Petitjean et al. (2000) concluded that the non-detection of molecular hydrogen in most of the DLA systems could be a direct consequence of high kinetic temperatures, $T>3000$ K, implying low formation rates of H$_2$ onto dust grains. Therefore, most of the DLA systems probably arise in warm and diffuse neutral gas. Fig.~5 of Petitjean et al. (2000) also suggests that, even if the gas is warm, H$_2$ should invariably be detected with $f\\la 10^{-6}$. Such very low molecular fractions probably result from high ambient UV flux.  In this paper, we present new results obtained from VLT-UVES high resolution spectroscopy of a DLA system at $z_{\\rm abs}=1.962$ toward Q\\,0551$-$366. Details of the observations are given in Sect.~2. The metal content, dust depletion pattern and H$_2$ molecular content of the absorber are discussed in, respectively, Sects.~3, 4 and 5. Our results are finally summarized and their implications discussed in Sect.~6. ",
        "conclusions": "\\begin{table*} \\caption{Physical conditions in individual components} \\begin{tabular}{cccccccc} \\hline \\hline $z_{\\rm abs}$ & $\\log N($C\\,{\\sc i}$)$ & $\\log N($C\\,{\\sc i}\\,$^\\star)$ & $\\log N($C\\,{\\sc i}\\,$^{\\star\\star })$ & $b$           & $T_{\\rm CMB}$ & $T_{\\rm 01}$     & $n_{\\rm H}$\\\\ &                        &                                & & (km s$^{-1}$) & (K)           & (K)              & (cm$^{-3}$)\\\\ \\hline 1.96152 & $12.69\\pm 0.07$  & $<12.18$\\phantom{\\,:} & $<11.94$\\phantom{\\,:} & \\phantom{0}$4.3\\pm 1.4$ & $\\le 10.5$ &   ...       & $<17$    \\\\ 1.96168 & $12.64\\pm 0.07$  & $12.84\\pm 0.07$       & $<12.16$\\,:           & \\phantom{0}$2.1\\pm 0.7$ & $\\le 22.3$ & $76\\pm \\phantom{0}7$   & $170-185$\\\\ 1.96180 & $12.42\\pm 0.13$  & $<12.18$\\phantom{\\,:} & $<11.94$\\phantom{\\,:} & \\phantom{0}$3.9\\pm 2.3$ & $\\le 12.9$ &   ...       & $<56$    \\\\ 1.96214 & $12.66\\pm 0.12$  & $12.69\\pm 0.11$       & $12.11\\pm 0.34$       & \\phantom{0}$2.1\\pm 0.8$ & $\\le 18.7$ & $81\\pm 29$  & $\\phantom{0}55-390$ \\\\ 1.96221 & $13.16\\pm 0.06$  & $12.98\\pm 0.09$       & $12.26\\pm 0.36$       &           $12.8\\pm 1.7$ & $\\le 14.5$ & $97\\pm 40$  & $\\phantom{0}30-150$ \\\\ 1.96268 & $12.63\\pm 0.08$  & $<12.18$\\phantom{\\,:} & $<11.94$\\phantom{\\,:} & \\phantom{0}$4.0\\pm 1.8$ & $\\le 10.0$ &   ...       & $<25$    \\\\ \\hline \\hline \\end{tabular} \\label{tabphy} \\end{table*}   Considering the cosmic microwave background radiation (CMBR) to be the only source of C\\,{\\sc i} excitation (e.g. Srianand et al. 2000), we can derive upper limits on the CMBR temperature from the column densities of C\\,{\\sc i}, C\\,{\\sc i}\\,$^\\star$ and C\\,{\\sc i}\\,$^{\\star\\star }$ that we measure in different components of the DLA system toward Q\\,0551$-$366 (see Table~\\ref{tabphy}). In the components where H$_2$ is detected, the values are found to be much larger than what is predicted from standard Big-Bang cosmology (i.e. about 8 K). Under the conditions prevailing in these clouds, fluorescence is negligible in populating the excited levels of C\\,{\\sc i} (see Fig.~2 of Silva \\& Viegas 2002). This means that excitation by collisions is important and, therefore, that the influence of local physical conditions (both density and temperature) is substantial in those clouds. Conversely, assuming the CMBR temperature to be 8 K and the kinetic temperature being approximated by the excitation temperature of the $J=1$ rotational level, we can estimate the hydrogen density from the relative populations of the C\\,{\\sc i} fine-structure levels. For the components where H$_2$ is not detected, we assume $T_{\\rm kin}=100$ K. As can be seen from Table~\\ref{tabphy}, it is apparent that the densities are larger in the gas where H$_2$ is detected. Therefore, the non-detection of H$_2$ in components having {\\it similar dust-to-metal ratios and similar column densities} of heavy elements is a consequence of lower densities and possibly higher temperatures. This is in line with the conclusions of Petitjean et al. (2000). The allowed range in pressure $p/k$ measured in the three H$_2$ components are respectively 12950$-$13770, 6050$-$20280 and 4170$-$8550 cm$^{-3}$ K. Such high pressures are seen only in $\\le 3$\\% of the C\\,{\\sc i} gas of our Galaxy (Jenkins \\& Tripp 2001) but are consistent with the high pressures also derived in the case of Q\\,0013$-$004 (Petitjean et al. 2002). They probably indicate that the gas under consideration is undergoing high compression due to either collapse, merging and/or supernovae explosions.   It is difficult to measure the rate of H$_2$ formation onto dust grains in DLA systems because it is related to the amount of dust which is itself difficult to estimate. In the present DLA system, even if the mean metallicity is close to solar the observed mean depletion of Fe is an order of magnitude smaller than what is typically measured in Galactic cold disc clouds. Therefore, we can estimate that the rate of H$_2$ formation onto dust grains is about one tenth of the Galactic value, thus $R\\sim 3\\times 10^{-18}$ cm$^3$ s$^{-1}$. Using the simple model proposed by Jura (1975) and the column densities and densities estimated above, we derive the photo-absorption rate of the Lyman and Werner bands of H$_2$: $\\beta_0=1.4\\times 10^{-11}-2.8\\times 10^{-10}$ s$^{-1}$ for the component at $z_{\\rm abs}=1.96214$. This value is similar to the average value measured in the Galactic ISM, $\\beta_0=5\\times 10^{-10}$ s$^{-1}$, while the average for the intergalactic radiation field is $\\beta_0\\approx 2\\times 10^{-12}$ s$^{-1}$ with $J_{21}({\\rm 912~\\AA })=1$. This shows that there is in-situ star-formation close to the absorbing clouds and that H$_2$ can survive, at least with such molecular fractions, even in the presence of a strong UV radiation field because the gas has a high particle density. One should keep in mind that the rate of H$_2$ formation goes linearly with the density of dust grains while it goes as the second power of the H\\,{\\sc i} density. Therefore, it is natural that, even though the presence of dust is a crucial factor for the formation of molecules, {\\it local physical conditions} such as particle density, temperature and local UV field, {\\it play an important role in governing the molecular fraction of a given cloud} in DLA systems."
    },
    "0206/astro-ph0206149_arXiv.txt": {
        "abstract": "s{The pulse profiles  of the accreting X-ray millisecond  pulsar \\sax\\ at different energies   are studied. The two main emission component, a black body and a power-law tail, clearly identified in the time-averaged spectrum, do not vary in phase. We show that the observed variability can be easily explained if the emission patterns of the black body and the Comptonized radiation are different: a ``knife\" and a ``fan\"-like, respectively. We suggest that Comptonization in  a hot slab (radiative shock) of Thomson optical depth $\\sim$ 0.3 at the surface of the neutron star may be responsible for the emission. We construct a detailed model of the X-ray production accounting  for the Doppler boosting, relativistic  aberration and gravitational light bending. The model reproduces well the  pulse profiles at different energies simultaneously, corresponding phase lags, as well as the time-averaged spectrum. By fitting the observed pulse profiles we obtain constraints on the neutron star radius ($R=7.5\\pm1.0$ km), the inclination of the system $i>60\\deg$, and the angle between the magnetic dipole and the rotational axis $\\delta=20\\deg\\pm 5\\deg$. } ",
        "introduction": "The X-ray burster \\sax\\ was detected in September 1996 using the Wide Field Camera aboard {\\it BeppoSAX} (in't Zand et al. 1998). Coherent pulsation with a period of 2.49 ms were detected in the {\\it Rossi X-ray Timing Explorer} data of the April 1998 outburst of this source (Wijnands \\& van der Klis 1998). Chakrabarty \\& Morgan (1998)  have discovered changes in the pulse arrival corresponding to the binary period of 2 hours. During the 1998 outburst, the source energy spectrum did not change much and the coherent pulsations with rms amplitude of 5-7 per cent were detected in all but one observation (Gilfanov et al. 1998, Cui et al. 1998) in spite of the 100-fold decrease of the luminosity. The pulse profiles show strong energy dependence and soft phase lags (Cui et al. 1998). By analyzing the  phase-resolved spectra, Gierli\\'nski et al. (2002) showed that the two main spectral components (soft black body and hard Comptonized) do not vary in phase. The black body emission is lagging the Comptonized emission. This causes changes in the pulse profile with energy and produces the observed soft phase  lags. While the black body light curve is nearly sinusoidal, the hard component is clearly asymmetric and cannot be fitted by a single harmonic. Gierli\\'nski et al. suggested that the Doppler effect can play a role in changing the pulse profile at high energies.   The source was observed 12 times by the \\xte\\ during the April 1998 outburst. We construct the energy-dependent pulse profiles following the procedure described in Gierli\\'nski et al. (2002) correcting the photon arrival times for orbital motions of the pulsar and the spacecraft. In order to improve the statistics, we average the profiles using the data from 11--26 April 1998.  In this paper, we construct a detailed model for the X-ray emission from a hot spot at the surface of the neutron star accounting for relativistic effects. We compare the model with the data and constrain the model parameters. ",
        "conclusions": "The two main spectral components (black body and a Comptonized) are clearly identified in the time-averaged spectrum of \\sax.  The pulse profiles show significant energy dependence and the softer photons are lagging the harder ones. Such a behaviour can be easily reproduced by a model where hard Comptonized emission peaks earlier than the black body emission (Gierli\\'nski et al. 2002). Due to the high velocity of the emitting region, the Doppler effect strongly affects the shape of the profiles. A black body  component shows about 25\\% variability amplitude and the Doppler effect (producing a $\\sim$ 6\\%  variability) does not change it much. We showed that  the Comptonized emission from a slab of Thomson optical depth of $\\tau\\sim 0.3$ has a broader angular distribution (a ``fan''-like) and therefore an intrinsically smaller variability amplitude. The  Doppler effect then shifts significantly the peak of the pulse at high energies where the Comptonized emission dominates producing soft lags.   By fitting the observed pulse profiles we obtain constraints on the neutron star radius ($R=7.5\\pm1.0$ km) and the inclination of the system $i>60\\deg$. We show that the angle between the magnetic dipole and the rotational axis is $\\sim 20\\deg$."
    },
    "0206/astro-ph0206239_arXiv.txt": {
        "abstract": "We investigate possibilities of solar coronal heating by acoustic waves generated not at the photosphere but in the corona, aiming at heating in the mid- to low-latitude corona where the low-speed wind is expected to come from. Acoustic waves of period $\\tau\\sim 100$s are triggered by chromospheric reconnection, one model of small scale magnetic reconnection events recently proposed by Sturrock. These waves having a finite amplitude eventually form shocks to shape sawtooth waves (N-waves), and directly heat the surrounding corona by dissipation of their wave energy. Outward propagation of the N-waves is treated based on the weak shock theory, so that the heating rate can be evaluated consistently with physical properties of the background coronal plasma without setting a dissipation length in an ad hoc manner. We construct coronal structures from the upper chromosphere to the outside of 1AU for various inputs of the acoustic waves having a range of energy flux of $F_{\\rm w,0}=(1-20)\\times 10^5$erg cm$^{-2}$ s$^{-1}$ and a period of $\\tau=60-300$s. The heating by the N-wave dissipation effectively works in the inner corona and we find that the waves of $F_{\\rm w,0}\\ge 2\\times 10^5$erg cm$^{-2}$s$^{-1}$ and $\\tau \\ge 60$s could maintain peak coronal temperature, $T_{\\rm max} > 10^6$K. The model could also reproduce the density profile observed in the streamer region. However, due to its short dissipation length, the location of $T_{\\rm max}$ is closer to the surface than the observation, and the resultant flow velocity of the solar wind is lower than the observed profile of the low-speed wind. The cooperations with other heating and acceleration sources with the larger dissipation length are inevitable to reproduce the real solar corona. ",
        "introduction": "The heating of the solar is been still poorly understood, and it is one of the most challenging but interesting questions to be solved in astrophysics. The origin of energy that heats the corona is generally believed to lurk in the turbulent convective motions beneath the photosphere. A certain fraction of the kinetic energy of those turbulent motions is carried up to the corona in the shape of non-thermal energy, such as magnetic or wave energy, and thermalization of such energy in the corona results in the heating of the surrounding plasma. Granule motions of the surface convection simply expect wave generation at the photospheric level. Possibilities of coronal heating by those waves have been investigated by many researchers (Osterbrock 1961; Ulmschneider 1971; McWhirter, Thonemann, \\& Wilson 1975 ). To date, among various modes of waves, although an Alfv\\'{e}n wave has widely taken up as a convincing candidate in coronal heating and acceleration of the high-speed wind in the polar coronal holes through the ion-cyclotron damping mechanism (e.g. Cranmer, Field, \\& Kohl 1999; Hollweg 1999), acoustic waves have not been regarded as a major heating source of the corona because of their dissipative character. Acoustic waves with a finite amplitude inevitably steepen their wave fronts to form shocks when they travel. They are strongly damped through upward propagation in the chromosphere, consequently the only 0.01\\% of the initial wave energy could reach the corona (Stein \\& Schwartz 1972; hereafter SS). A role of acoustic waves in solar coronal heating has been buried in oblivion for a long time.  However, recent observations of dynamical structures of the solar corona (see Aschwanden, Poland, \\& Rabin 2001 for recent review) highlight a role of acoustic waves generated not at the photosphere but in the corona. A large number of flares and flare-like events (e.g. Tsuneta et al.1992; Tsuneta 1996) have been observed by various telescopes until today (e.g. Nishio et al.1997; see Bastian, Benz, \\& Gary 1998 for review) in a wide range of energy from $\\sim 10^{24}$erg to $\\sim 10^{32}$erg, and they are statistically argued in terms of supply of bulk energy to heat the global corona. Recent measurements of EUV frequency distribution, $\\frac{dN}{dE}(E) \\propto E^{-\\alpha}$, of flare-like events gives a steeper power-law index ($\\alpha=2.3 - 2.6$ by Krucker \\& Benz 1998; $\\alpha=2.0 - 2.6$ by Parnell \\& Jupp 2000) on the lower energy side ($E=10^{24}-10^{25}$erg). This fact makes us infer that small-scale flare-like events, such as nano-flares, might be sufficient (Hudson 1991) for the required energy budget ($\\gtrsim 10^{5.5}$erg cm$^{-2}$s$^{-1}$; Withbroe \\& Noyes 1977). Several models taking into account such small-scale events have been introduced (Sturrock 1999; hereafter S99; Sturrock, Roald, \\& Wolfson 2000; Roald, Sturrock, \\& Wolfson 2000; hereafter RSW; Tarbell, Ryutova, \\& Covington 1999; Sakai et al. 2000; Furusawa \\& Sakai 2000). Among these models, S99 proposed a model of chromospheric magnetic reconnection. In this model, a flux tube, newly formed as a result of reconnection events, oscillates vertically to excite acoustic waves in the corona. RSW further showed based on a simple kinetic model that this process possibly liberates the sufficient energy in the corona with the correct relevant parameters. A different mechanism also predicts generation of acoustic waves in the corona. For example, spicules in the open magnetic field region effectively transports energy of random motions at the photosphere into the corona and produce longitudinal (acoustic) waves there \\citep{hol92,ks99}. Thus, acoustic waves are supposed to be generated constantly in the corona by various dynamical processes. Such acoustic waves could heat the surrounding corona directly, unlike acoustic waves created at the photospheric level. They are expected to heat the inner corona effectively by the shock dissipation. Therefore, density at the coronal base might be large enough to creat dense wind flow observed as the streamer in the mid- to low-latitude region, which is believed to be connected to the low-speed solar wind \\citep{hwf97}. Acoustic waves might become one of relevant processes working in the region generating the low-speed wind of which mechanisms have been poorly elucidated.   In the above models introducing production of acoustic waves in the corona, the authors concentrated on the amount of kinetic (wave) energy released in the corona. However, heating of coronal gas is accomplished by thermalization of wave energy through dissipation of such acoustic waves. Therefore, an appropriate treatment of propagation and dissipation of these waves is indispensable in understanding a problem of coronal heating. In this paper, we employ a formulation of a weak-shock theory constructed by SS, originally developed to study acoustic waves generated by convective motions at the photosphere, for propagation and dissipation of the waves produced in the corona. Then, we can explicitly determine the rate of heating by dissipation of the waves on given wave energy flux (or amplitude) and wave period without tuning other free parameters; we do not have to arbitrarily take an exponential type of heating of mechanical energy input, $F_{\\rm m} \\propto \\exp(\\frac{-(r-R_{\\odot})}{l_{\\rm m}})$, for an assumed constant dissipation length, $l_{\\rm m}$, which is poorly supported by fundamental physical processes, although this conventional shape of the heating law was adopted, for lack of an alternative in most previous models \\citep{ko76,wtb88,sl94}, to study the global coronal structure.  The paper is organized as follows: In \\S\\ref{sc:mdl}, we present our model. We briefly summarize the generation of acoustic waves in the corona (\\S\\ref{sc:cmr}), and the weak-shock theory formulated by SS (\\S\\ref{sc:wst}). Basic equations describing the global coronal structure from $1 R_{\\odot}$ to $\\gtrsim$ 1AU are shown in \\S\\ref{sc:be}, and we present our method to construct a unique coronal solution on given wave parameters. In \\S\\ref{sc:rslt} we show our model results. First, we discuss variation of coronal wind structures with respect to input parameters from a theoretical point of view (\\S\\ref{sc:awdp}). Second, we examine several characteristic properties of the resultant corona (\\S\\ref{sc:rscp}). Finally, we test whether our model is able to reproduce recent observed results for the low-speed wind in the mid- to low-latitude region (\\S\\ref{sc:obs}). In \\S\\ref{sc:smds}, we summarize our results and discuss related topics. ",
        "conclusions": "\\label{sc:smds} We have investigated the process of acoustic waves generated in the corona as a heating source especially in the mid- to low-latitude corona where the low-speed winds come from. We have found that the acoustic waves with $\\tau\\ge 60$s and $F_{\\rm w,0} \\ge 2\\times 10^5$erg cm$^{-2}$s$^{-1}$ could certainly heat up the ambient plasma to $T\\ge 10^6$K by the dissipation of the N-waves even in a flow tube with the areal expansion of $f_{\\rm max}=5$, by balancing with the losses of the radiative cooling, the downward thermal conduction, and the adiabatic loss by the solar wind. Due to its dissipative character, the dissipation of N-waves effectively works in the inner corona and reproduces the density profile observed in the streamer region. However, it cannot contribute to the heating of the outer corona, since most of the wave energy is damped within a region of few-tenth of the solar radius even if one considers the waves with a long period of $\\tau=300$s. As a result, it is impossible to explain the observed temperature profile and flow velocity of the low-speed wind only by that process. Therefore, other mechanisms with the larger dissipation length should play a role cooperatively in the coronal heating and the acceleration of the low-speed wind.  Small-scale reconnection events and spicules also predict generation of fast shock waves \\citep{lw00,lee01,hol82}, though we have concentrated on the role of the N-waves (slow shocks along the magnetic field line) excited by those events in this paper. The dissipation length of the fast shocks must be larger than N-waves with the identical period, since their phase speed is $\\gtrsim v_{\\rm A}$, which is much larger than that of the N-waves ($\\sim c_{\\rm s}$) for low-$\\beta$ coronal plasma. Hence, cooperation with the fast shocks would give a more distant location of $T_{\\rm max}$ to match the observed temperature profile (fig.\\ref{fig:obste}), and would let the acceleration of the wind flow continue in the outer corona, as suggested by observation of the streamer belt (fig.\\ref{fig:obsvl}).  Finally, we had better remark an issue on anisotropic and selective heating of ions. According to recent observations \\citep{ssp02}, O VI ions in the streamer have high perpendicular kinetic temperature, though this is not so extreme as that observed in the high-speed wind. Our results show that the deposition of energy and momentum from the N-waves to the ambient gas is completed in the region of $r\\lesssim 1.5 R_{\\odot}$, where the particles are thermally well coupled. As a result, whatever anisotropies obtained by the N-wave heating would be wiped out and all the ions would have isotropic kinetic temperature. The process of the N-waves cannot explain the observed anisotropies, which also indicates that the other mechanisms causing the anisotropic heating have to operate simultaneously.   We thank K. Shibata, K. Ohki, K. Omukai, K. Tomisaka, H. Saio, S. Nitta, S. Inutsuka, T. Kudoh, Y. Yoshii, T. Kajino, and A. Tohsaki as well as members of DTAP in NAOJ for many valuable and critical comments and Y. Mclean for improvement of presentation in this paper. The author is supported by the JSPS Research Fellowship for Young Scientists, grant 5936."
    },
    "0206/hep-th0206149_arXiv.txt": {
        "abstract": "As a simple model for unknown Planck scale physics, we assume that the quantum modes responsible for producing primordial curvature perturbations during inflation are placed in their instantaneous adiabatic vacuum when their proper momentum reaches a fixed high energy scale $M$. The resulting power spectrum is derived and presented in a form that exhibits the amplitude and frequency of the superimposed oscillations in terms of $H/M$ and the slow roll parameter $\\epsilon$. The amplitude of the oscillations  is proportional to the third power of $H/M$.  We argue that these small oscillations give the lower bound of the modifications of the power spectrum if the notion of free mode propagation ceases to exist above the critical energy scale $M$. ",
        "introduction": "Inflation, in addition to explaining the homogeneity and flatness of the universe, provides a platform for computing the properties of the small  inhomogeneities which give rise to the large scale  structures \\cite{L90,LL00}.  When initially in the vacuum, the quantized fluctuations of the coupled system of metric and scalar inflaton field produce a nearly scale invariant spectrum. Small deviations arise from the time dependence of the  Hubble expansion rate during inflation \\cite{MFB92}.  The  predictions of the simplest inflationary scenarios are in good agreement with recent analysis of high-resolution CMBR maps.   However, these predictions are based on the incomplete framework of semiclassical quantum field theory.  One particular aspect which has been questioned lately is the choice of the ground state of the quantum modes that generate the fluctuations. In the standard derivation, the vacuum is imposed at  the beginning of inflation for all quantum modes. During inflation,  the proper wavelengths of these modes are redshifted by many orders of magnitude.  Given that inflation lasts more than 70 e-folds, the initial state of the modes responsible for the cosmological structures  observed today is assigned at proper wavelengths shorter than the Planck length. Moreover, in the semiclassical framework, one implicitly assumes that their redshift to lower  momenta proceeds without interference from Planck scale physics.  The validity of these assumptions is unclear; this is known as the trans-Planckian question of inflationary cosmology \\cite{J00}.  The sensitivity of the predictions of inflation, in particular the scale invariance of the power spectrum, with respect to trans-Planckian physics has been probed in a number of different ways, including nonlinear dispersion relations \\cite{MB00,BM00,BJM01,N00,NP01,S02}, string-inspired short  distance cutoffs \\cite{K00,KN01,Eea01,Eea01b,HS02}, and modifications derived from non-commutative geometry \\cite{CGS00,Lea02b,BH02}. Two aspects of the discussion have become particularly  important: the question of the leading order correction to the power spectrum in terms of the ratio (called $\\sigma$ in this work) of the Hubble parameter $H$ and the scale of new physics $M$, and the existence of distinctive features of short distance effects that would allow their unambiguous detection in a general inflationary setting. As a candidate for the latter, the potential violation of the scalar-tensor consistency relation by Planck scale physics has been mentioned \\cite{HK01}. Unfortunately it requires the measurement of the tensor perturbation power spectrum with great accuracy \\cite{Kea02}. Alternatively, if models for short distance physics predict a unique pattern imprinted on the scalar power spectrum, this alone might be sufficient for detection.  With regard to the amplitude of the correction, Ref.\\cite{Kea02} has pointed out that local  effective field theory generically allows only $\\sigma^2$ and smaller effects, whereas \\cite{BM02} have argued that $O(\\sigma)$ effects are possible if adiabaticity of the mode evolution is violated.  Danielsson  \\cite{D02} has indeed found a linear correction by imposing  that each quantum mode obeys a positive frequency solution at a given energy scale $M$ (see Eq.(33) in  \\cite{Eea02}). This method has been re-analyzed in Ref.\\cite{Eea02} in a general inflationary background, confirming the $O(\\sigma)$-amplitude and demonstrating the oscillatory character of the correction to the power spectrum. It has also been noted that some classes of cutoff models predict corrections that depend on $M$ in more complicated ways \\cite{SW02}.  In this work, we follow Danielsson's approach but we replace his vacuum condition  by the instantaneous adiabatic vacuum at the scale $M$.  The basic underlying physical assumption is that whatever physics determines the dynamics at higher energies, it places the quantum modes into their adiabatic vacuum at  some scale $M \\lesssim m_{\\rm Planck}$, from where on they evolve freely and in the usual way, i.e. without dispersion.  Owing to space-time expansion, this state is not exactly identical to the asymptotic vacuum which is imposed in the infinite past. This difference gives rise to a modification of the power spectrum whose shape and amplitude we calculate both exactly and perturbatively. In contrast with \\cite{D02,Eea02}, we find a leading correction that is cubic in $\\sigma$, thus predicting a much smaller modification of the spectrum.  The procedure of assigning the vacuum when the proper momentum of each mode crosses $M$ can be considered a minimal model for short distance effects since, in an expanding space-time, the adiabatic vacuum is the closest analogue to the usual Minkowski vacuum.  Indeed, any other choice of the initial state will  produce larger effects  unless the amplitudes to find pairs of adiabatic quanta are fine tuned so as to ``undo'' the Bogoliubov transformation with respect to the asymptotic  vacuum, see the conclusions for more details.  Hence the modification of the power spectrum we found  can be interpreted as the minimal one if modes are ``created'' at the scale $M$.  Regarding the pattern that short distance physics imprints on the microwave background, Easther et al. \\cite{Eea01b} have shown numerically for one specific model that nearly sinusoidal oscillations in $\\log(k)$ may be superimposed on the power spectrum. This has been confirmed in \\cite{Eea02} using the cutoff criterion of \\cite{D02}. We also find a distinctive oscillatory pattern and present a simple expression for its oscillations in wavenumber space. These oscillations arise through the wavenumber dependence of $\\sigma$, which is governed by the parameter $\\epsilon$ measuring the deviation of the inflationary background from de Sitter space. Since in the slow-roll framework, $\\epsilon$ can be determined from the spectral index and the tensor/scalar ratio \\cite{K98b}, our model makes an explicit and simple prediction for the amplitude and shape of the correction. We hope that this result will be useful in future studies of the detectability of short distance effects in cosmological measurements. ",
        "conclusions": "In spite of the smallness of $\\delta$ in our model, this result can be seen in an optimistic light. $\\sigma$ may be larger than the   normally assumed value of $10^{-5}$, allowing a potential detection of short distance effects in CMBR measurements \\cite{Eea01b,Kea02}, especially in combination with a measurement of the tensor spectrum. Moreover, it is by no means obvious that nature chooses to place newly born modes in the instantaneous adiabatic vacuum at $\\eta_k$.  Indeed, when using a nonlinear dispersion relation which modifies $\\omega_k(\\eta)$ for $k/a > M$, one might obtain a state characterized by Bogoliubov coefficients $\\alpha_k, \\beta_k$ which signal the departure from the adiabatic vacuum. In this case, the leading correction will be given by $\\vert \\delta_k \\vert \\simeq \\vert  \\beta_k/\\alpha_k + \\beta^{\\eta_k}_k/\\alpha^{\\eta_k}_k\\vert$ in place of Eq. (\\ref{delta}). At this point, two remarks should be made.  First, since $\\vert \\delta_k \\vert $ is given by a sum of amplitudes and not a sum of squares, one cannot exclude the possibility of a partial cancellation of the two Bogoliubov transformations. However, an exact cancellation is unlikely since $\\beta_k$ and $\\beta_k^{\\eta_k}$ have different origins.  Secondly, one should keep in mind that the coefficients $\\beta_k$  must obey $\\vert \\beta_k \\vert \\lesssim \\sigma_k$  in order not to  violate the hypothesis that the dominant contribution  to the Friedman equation be the inflaton potential energy \\cite{T00,NP01,S01}. Similar constraints follow from the particle production rate in the late universe and limits on the observed gamma-ray background \\cite{ST02}. Whether there exists a window in parameter space that allows the detection of a high-energy cutoff in the CMBR remains an interesting question that, in our opinion, cannot be fully answered at this point.  The simple prediction for the frequency of the superimposed oscillation in the power spectrum given by Eq.(\\ref{coeffs}) may allow a systematic search for such features in cosmological data. If an oscillation is detected whose frequency, for some $\\sigma$, is consistent with the value of $\\epsilon$ inferred from the scalar spectral index and the tensor/scalar ratio \\cite{K98b}, it would represent a strong indication of modified short distance physics. It is certainly worthwhile exploring the parameter space that would make this detection feasible."
    },
    "0206/astro-ph0206262_arXiv.txt": {
        "abstract": "We perform an accurate computation of the production rate for primordial black holes (PBHs). The reason is that the underlying mass variance had been overestimated systematically, as was shown recently. For scale-free powerlaw primordial spectra, and for a Universe with critical density, the mass variance is less than 34\\% of its value thought earlier for the spectral index in the range $1\\leq n\\leq 1.3$. We then extend our study to spectra with a characteristic scale and find the accurate shape of the corresponding mass variance. For a pure step in the primordial spectrum, the step in the variance is smoothed around the characteristic scale $k_s$. For a spectrum with large oscillations near $k_s$, we find a pronounced bump in the variance. This could yield a significant part of the cold dark matter in the form of PBHs with mass $M$ in the range $5\\times 10^{15}\\ {\\rm g}\\lesssim M \\lesssim 10^{21}\\ {\\rm g}$. ",
        "introduction": "Primordial black holes (PBHs) are black holes that result from the collapse of density fluctuations in the early Universe \\cite{CH74}. This is the only way to produce black holes with mass smaller than about three solar masses (except possibly by the production of black holes in accelerators, as has recently been discussed in the speculative framework of branes). This makes them interesting for several reasons. First, PBHs with initial mass $\\lesssim M_*\\approx 5\\times 10^{14}\\ {\\rm g}$ have been already evaporated by Hawking radiation or evaporate at the present stage of the Universe. They would play a key role in studying quantum gravitational effects of black holes (see, e.g., \\cite{GBH}). Second, PBHs with any mass could be an excellent indicator of the conditions in the early Universe. And third, they could in principle contribute a major part to dark matter. As no PBHs have yet been observed (or identified as such), existing observational limits can strongly constrain scenarios of the early Universe.  One of the most successful scenarios is inflation (see, e.g., \\cite{LL} for a review). It leads to a primordial spectrum of fluctuations that can serve as the seed for the observed structure in the Universe. Recent observations of the anisotropies in the cosmic microwave background (CMB) radiation strongly support inflation. Depending on the exact form of the spectrum, the same fluctuations can also lead to the formation of PBHs (see, e.g., \\cite{Carr2,GL}). Most models of inflation predict a scale-free spectrum with a spectral index $n$ very close to the scale-invariant case $n=1$. A significant number of PBHs can only be produced for $n>1$ (a ``blue spectrum''), since these values lead to more power on small scales (this was first addressed in \\cite{CGL}). Observational limits (both from Hawking radiation and the fact that PBHs must not overclose the Universe) strongly constrain $n$. This, therefore, yields a constraint on inflationary models that is independent of the cosmological constraints from the CMB and large-scale structure.  It has been found recently that the production rate of PBHs was systematically overestimated \\cite{BKP1}. The reason is that one has to use the transfer function (the tool to describe the evolution of fluctuations within the horizon) and evaluate it at the onset of PBH formation, not at the present time. The qualitative explanation goes as follows: at each stage of the cosmological evolution, the correct use of the transfer function takes into account the dynamics of the fluctuations after horizon-entry. This leads effectively to very different spectra on small scales at different times. A general formula to remedy this earlier deficiency has been presented in \\cite{BKP1} and some quantitative estimates have been given. The extension from vanishing cosmological constant discussed in \\cite{BKP1} to the nonvanishing case has been performed in \\cite{P02}. However, an accurate calculation has to be done numerically. This will be done here and will constitute the first aim of the present paper.  Although scale-free models have been successful so far, it is not clear that the true spectrum must be exactly scale-free. The inflaton potential may exhibit distinguished features at an energy scale relevant to GUT theories. This would lead to models with broken scale invariance (BSI). One simple example would be a jump in the derivative of the inflaton potential, interrupting the slow-roll behavior and leading to additional particle creation and formation of PBHs. The spectrum of fluctuations resulting from such a jump can be presented analytically in closed form \\cite{S92}. It can cope surprisingly well with observational constraints like the possible presence of a bump in the matter power spectrum \\cite{LPS98}. The BSI model contains two additional parameters: the strength of the jump, $p$, and its location, $k_s$.  In \\cite{BKP1} we have already discussed the simple case of a pure step in the mass variance. The BSI model of \\cite{S92} is, of course, of much more physical relevance. The second aim of our paper is thus to perform an accurate calculation of PBH production in this model. In fact, the production rate can be so high that the PBHs could be candidates for massive compact halo objects (MACHOs) and thus contribute a significant part of the dark matter in the Universe \\cite{BKP2}.  BSI models have the advantage that they are the only realistic models of PBH formation. The reason is that scale-free models need a certain fine-tuning to lead to a significant number of PBHs: if the spectral index is only a little more than some critical value $n^*$, the PBH production rate would be so high that it were in conflict with observation. On the other hand, for $n$ being somewhat smaller than $n^*$, the production rate would be exponentially small and one would never have the chance to observe a PBH. The critical value is $n^*\\approx 1.3$ \\cite{BKP1} which is slightly bigger than current observational limits (see e.g. \\cite{Max01}). Even more, a blue spectrum with constant spectral index close to $n^*$ is inevitably ruled out due to the overproduction of PBHs which have evaporated by now and contribute to the diffuse $\\gamma$-ray background. The situation will be drastically different in a BSI model, as will become clear from the results presented here. PBH production from models with a scale was considered before \\cite{INN94,RSG96,BP}, but we shall reconsider it here in the light of our improved formulae from \\cite{BKP1}.  Our paper is organized as follows. In Sec.~II we shall briefly review the fundamental formalism to calculate PBH formation, see \\cite{CH74} and \\cite{BKP1} for more details. In particular, we shall present the crucial relation between primordial power spectrum and mass variance. In Sec.~III we then give accurate results for the cases of scale-free spectra, spectra with a pure step, and the BSI model of \\cite{S92}. Sec.~IV is devoted to the discussion of observational constraints for each of these situations. For the BSI model, in particular, these constraints allow the possibility to form PBHs in the mass range $5\\times 10^{15}\\ {\\rm g}\\lesssim M \\lesssim 10^{21}\\ {\\rm g}$. Such PBHs could play the role of MACHOs in the halos of galaxies and could contribute significantly to the dark matter in the Universe. ",
        "conclusions": "New expressions were derived recently for the improved computation of the mass variance at horizon crossing \\cite{BKP1}. We had found that for a scale-free powerlaw primordial spectrum, the mass variance for PBH masses around $M_*\\approx 5\\times 10^{14}\\ {\\rm g}$ is significantly lower than it had been assumed in the literature. In this paper we have refined our estimate and confirmed this discrepancy: We find that $5.37\\leq \\alpha^2(M_*)\\leq 6.83$, for $1\\leq n\\leq 1.3$, see Eq.~(\\ref{alphanum}), in particular $\\sigma^2_H(t_k)= 6.63~\\delta^2_H(t_0)$ for a scale-invariant (Harrison-Zel'dovich) primordial spectrum. For scale-free primordial spectra, the quantity $\\alpha$ is further independent of $M$ for a wide range of scales. This is to be compared with the value $\\sigma^2_H(t_k)\\approx 25 ~\\delta^2_H(t_0)$ which corresponds to the normalization adopted previously in the literature.  We have further applied our formalism to primordial spectra which are not scale-free. Two specific spectra possessing a characteristic scale were considered. The first one is purely phenomenological and corresponds to a pure step in the primordial spectrum. We find that the step is translated into the probability $\\beta(M)$ in a non-trivial way: It acquires a smoothed step centered around the characteristic mass $M_s$. The second primordial spectrum corresponds to a BSI spectrum which is obtained from inflationary perturbations when the inflaton potential has a jump in its first derivative at some scale $k_s$. It would of course be of interest to relate this scale to a scale of fundamental physics at the GUT scale. In this case, we find that the mass variance $\\sigma^2_H(t_k)$ displays a sharp, pronounced peak centered at the scale $k_{peak}\\approx 1.27 ~k_s$ which is {\\it not} the exact location of the maximal peak in the primordial spectrum itself. Again this is due to the non-trivial relationship between the primordial perturbations spectrum itself and the mass variance at Hubble radius crossing $\\sigma^2_H(t_k)$. This in turn yields a pronounced bump in the quantity $\\beta(M)$ around $M_{peak}$ which has interesting applications. We have considered here in detail one such application \\cite{BKP2}, the possibility to produce in this model a significant part of the cold dark matter in our Universe. Interestingly, we find that the mass of the produced PBHs is constrained to be much smaller than one solar mass. Indeed, the characteristic mass of the spectrum, essentially the same as that of the produced PBHs, cannot be fixed arbitrarily because of the severe constraint coming from the $\\gamma$-ray background. The inclusion of a non-negligible spectrum of primordial gravitational waves and running scalar and tensorial spectral indices $n_s$ and $n_T$ (see e.g.\\cite{DS01}) could enlarge the class of models that can realize our scenario. A thorough study should probably await a more careful determination of $\\delta_{min}$ and, to less extent, of the cosmological parameter $\\Omega_{\\Lambda}$. At the hand of the example studied here, it is clear that spectra with characteristic scales on very small scales, much smaller than those probed by the CMB anisotropy or large-scale structures, offer interesting new avenues.  We have studied here concrete models with concrete assumptions, for which we have obtained precise quantitative results. It is clear that these results may be different for other models. They may also depend on the assumptions made. In particular, the influence of non-Gaussian fluctuations could strongly change the PBH formation rate \\cite{BP}. However, all situations can in principle be dealt with using the formalism presented here."
    },
    "0206/astro-ph0206112_arXiv.txt": {
        "abstract": "{ The soft X-ray transient GU Mus has been detected by XMM-Newton in the quiescent state.  The source is very faint, with a 0.5-10.0 keV unabsorbed flux of $\\simeq 1.1 \\times 10^{-14}$ ergs cm$^{-2}$ s$^{-1}$.  The spectra is well fit by an absorbed powerlaw with a photon index of $ \\alpha = 1.6 \\pm 0.4$, close to the value seen when the source was in the low/hard state in Aug.  1991.  From our observed luminosity, it seems unlikely that the quiescent state emission is dominated by coronal X-rays from the secondary.  The flux also appears to be in agreement with the ADAF model of BH-transients in quiescence. ",
        "introduction": "Soft X--ray transients (SXTs) are compact binaries where a neutron star (NS) or black hole (BH) primary accretes from a Roche--lobe filling donor star.  While there exists unique observational signatures for the identification of a neutron star in these systems, e.g.  through X--ray bursts which are thought to be due to thermonuclear burning of superheated accreted matter on the surface of a neutron star, there is no similar identifier for an event horizon -- the unique characteristic of a BH-system.  A knowledge of the mass function of the system (from optical observations of the companion), can at most give a lower limit to the mass of the compact object.  In the case of stellar mass BHs, this is usually not very definitive, given the present uncertainty on the limiting mass of neutron stars.  The spectra of SXTs in outburst appear to be dominated by a radiatively efficient, optically thick, geometrically thin, accretion disc in a hot, high viscosity state.  However, spectra of BH-SXTs in quiescence do not appear to be a simple accretion disc phenomenon alone (e.g..\\ Lasota 2001).  A simple, two component model for SXTs in quiescence consists of an outer, thin, accretion disc, which serves as a reservoir for the accumulated mass.  An inner, hot, advection-dominated accretion flow (ADAF) transports energy from the inner edge of the accretion disc to the compact object (Menou et al.  1999).  In the case of NS-SXTs, this would heat up the surface of the NS (without thermonuclear runaway) causing it to radiate, or, in the case of BH systems, carry the energy, along with matter, out of sight beyond the event horizon.  Thus the quiescent state spectra of SXTs would yield information on the nature of the compact object.  Hence, in recent years the quiescent properties of SXTs and their white dwarf analogues (dwarf novae) have taken centre stage in the quest to prove the existence of black holes by detecting the event horizon, and in the search for the holy grail of accretion disc theory, the viscosity mechanism (Garcia et al.,2001, and Narayan et al., 2001).   The two-component model discussed above is not without inconsistencies. Inclusion of viscous heating of electrons, along with that of protons, accounting for convectively unstable flows, and magneto--rotational instability in 3-D numerical computations of ADAFs now show that ADAFs develop strong outflows (Hawley et al.\\ 2001).  Also, the large mass flow inherent in ADAF models predicts a much larger L$_{\\rm quiesc}$ in NS systems than observed. However, all ADAF models consistently predict that BH-SXTs should have a much lower quiescent flux (by a factor of $ \\simeq 100$) relative to the NS SXTs.  Prompted by the apparent failure of ADAF for NS-SXTs, alternative explanations for the quiescent emission have been put forward, including coronal emission from companion stars (e.g..\\ Bildsten \\& Rutledge 2000; but these exceed the existing quiescent BH SXT detections), and emission from the neutron star surface that was heated in earlier outbursts (Brown et al.\\ 1998).   So far, quiescent state detections, or upper limits, exist for 6 NS-SXTs and 8 BH-SXTs.  The quiescent luminosities of these systems in the 0.5-10.0 keV range (Table 2, Narayan et al., 2001), normalised to $10^{38}$ ergs s$^{-1}$ are plotted as a function of the orbital period $P_{\\rm orb}$ in fig.  \\ref{fig:  LvsP}. Systems with similar orbital periods are expected to be at the same point in their evolutionary history, and hence expected to have similar accretion rates $\\dot M$ (however, see Sect. \\ref{sec: Discussion}). In accordance with the ADAF models, the BH-SXTs do appear to be fainter than the NS SXTs.  We discuss the cause of this scatter in Sect. \\ref{sec:  Discussion}.  We had designed our observation of the BH-SXT GU Mus in order to (a) investigate the ADAF prediction of the relative faintness of quiescent BH systems over NS systems with similar orbital period and (b) to obtain a potential spectral signature for the existence of an event horizon -- i.e., the search for the indisputable evidence of a black hole.  GU Mus is an ideal candidate for testing these hypotheses because no X-ray activity has been detected from it since the last outburst in Jan.  1991.  We discuss the outburst properties, the present observation and our results in the following sections.  \\begin{figure}[h] \\resizebox{\\hsize}{!}{\\includegraphics{h3559_fig1.eps}} \\caption{Quiescent luminosities of BH-- and NS-- SXTs plotted against their orbital periods. The data is taken from Table 1 of Narayan et al. 2001. All luminosities are in the 0.5 to 10 keV band.} \\label{fig: LvsP} \\end{figure} ",
        "conclusions": "\\label{sec: Discussion}  Our observation had been planned with the intention of using the spectral characteristics of GU Mus to provide an observational discriminant between NS- and BH-SXTs.  However, we find that, while we have a clear detection of GU Mus in the quiescent state, at $\\sim 100 $ counts, the spectra have insufficient counts to constrain well even the simple absorbed power-law model.  By fixing the column density at $N_H = 0.26 \\times 10^{22}$ cm$^{-2}$, the fit to the EPIC-PN spectra constrains the photon index to $\\alpha = 1.6 \\pm 0.4$.  This appears to be consistent with both the outburst and post-outburst observations, where the photon index of the hard component decreased from its high/soft state value of 2.2-2.7 (Jan.  to Apr.  1991) to its low/hard value of $\\sim 1.6$ in Sept.  1991 (Ebisawa et al.  1994).  Besides GU Mus, only 5 other BH-SXTs have been actually detected in quiescence (for Chandra observations, see Kong et al. 2001, and Garcia et al.  2001) and these luminosities have been plotted in fig. \\ref{fig:  LvsP}, as a function of orbital period $P_{\\rm orb}$.  In addition, upper limits for the BH SXTs 4U 1543-47 ($P_{\\rm orb} = 27$ hr., Orosoz et al.  1998) and H1705-250 ($P_{\\rm orb} = 12.7 $ hr, Narayan et al, 1997) have also been plotted. The luminosities for the Chandra observations included in fig.  \\ref{fig:  LvsP}, were calculated (Narayan et al., 2001 and references therin) in the energy range 0.5-10.0 keV, by fitting to a power-law spectrum with photon index $\\alpha \\sim 2.0$ and column densities consistent with the optical extinction.   The evaluation of the intrinsic X-ray luminosity of GU Mus in quiescence is uncertain, chiefly because of the poor constraints on the distance.  In fig. \\ref{fig:  LvsP}, we have plotted the value of $\\log L_{quies}/10^{38} = -6.40 $ for GU Mus, using our reference value of 5.5 kpc. In terms of the Eddington luminosity, using an inclination of $i= 60^{\\circ}$, $L_{0.5-10.0}/L_{Edd} = 9.25 \\times 10^{-8} (d/5.5 \\rm kpc)^2$.  We plot the luminosities in fig.  \\ref{fig:  LvsP} as normalised to a factor of $10^{38}$ ergs s$^{-1}$, instead of the ratio $L_{\\rm quiesc}/L_{ \\rm Edd }$, as preferred by others (e.g.  Garcia et al 2001), because the Eddington luminosity $L_{ \\rm Edd }$ will be $ \\sim 10$ times larger in BH systems than in NS systems, thus leading to an artificial separation between the BH and NS systems.  A natural separation can be expected in compact binaries with P$\\leq 12$~hrs when the accretion rate $\\dot M $ is driven by orbital angular momentum losses $\\dot J$.  If these are dominated by magnetic braking, $\\dot J$ is the same for identical donor stars, so that $\\dot M $ is $\\propto 1/J$ i.e.\\ roughly to $1/M_{\\rm primary}$ (e.g.\\ King et al.\\ 1997). Hence $\\dot M$ in BH systems would be expected to be smaller than in NS systems. Normalisation relative to $L_{\\rm Edd}$ is reasonable only if the event were driven by gravitational radiation, since, in that case, $\\dot M$ would scale with $M$. However, the truth may be more complicated.  Some short--period NS systems may be remnants of a thermal--timescale mass transfer phase (King et al.\\ 2001), and at least one short--period BH system (XTE J1118+480) shows evidence nuclear--processed material (Haswell et al.\\ 2001), a clear sign of an evolved donor. The secondary in the BH-SXT GS 2000+25, too, has been found to be an evolved donor (Harlaftis et al, 1996).  All this suggests a greater spread of $\\dot M$ at a given $P_{\\rm orb}$.  An alternative explanation for SXT quiescent X-ray emission was given by Bildsten and Rutlidge (2000), who suggested that it arose from the corona of the (necessarily) rapidly rotating companion star.  In order to investigate this possibility, we fitted our data to the Raymond-Smith model for a coronal plasma, assuming solar abundances.  If the emission is indeed coronal in origin, rather than an ADAF in the inner disk region, then the temperature of the plasma should be $ \\le 1.4 $ keV (Dempsey et al.  1993), and the ratio of the soft X-ray to bolometric luminosities be $ L_{\\rm SX}/L_{ \\rm bol} < 10^{-3} $.  While our fits constrain the temperature too poorly to compare, we do note that the {\\it absorbed} flux in the 0.4-2.4 keV range is $3.3 \\times 10^{-15}$ ergs cm$^{-2}$ s$^{-1}$, (EPIC-PN spectra with constant $N_H$) which implies, $L_{\\rm  SX}/L_{ \\rm bol} > 8.96 \\times 10^{-3}$.  In order to calculate $L_{\\rm bol}$, we used $V = 20.5$ for GU Mus, with a $50 \\%$ contribution from the accretion disk (Orsoz et al. 1996), resulting in $V_{\\rm secondary} = 21.25$.  The other parameters were $E(B-V)=0.3$ ( Della Valle et al.  1991), $A_V = 3.1 E(B-V)$ and distance $d= 5.5$ kpc.  Thus, based on luminosity estimates alone, it maybe unlikely that the quiescent state X-ray emission originates from the companion star's corona.  We also note, however, that of all the other BH-SXTs observed in quiescence so far (GRO J0422+32, A0620-00, XTE J1550-564, V404 Cyg, GRO J1655-40 and GS 2000+25, Kong et al.  2000), the stellar corona appears to contribute a major portion of the soft X-ray luminosity only in GS 2000+25, although as yet no spectra exists for the quiescent state of this source."
    },
    "0206/astro-ph0206438_arXiv.txt": {
        "abstract": "We present $UBVRIz$ lightcurves of the Type~Ia SN~1999ee and the Type~Ib/c SN~1999ex, both located in the galaxy IC~5179. SN~1999ee has an extremely well sampled lightcurve spanning from 10 days before $B_{max}$ through 53 days after peak. Near maximum we find systematic differences $\\sim$0.05 mag in photometry measured with two different telescopes, even though the photometry is reduced to the same local standards around the supernova using the specific color terms for each instrumental system. We use models for our bandpasses and spectrophotometry of SN~1999ee to derive magnitude corrections (S-corrections) and remedy this problem. This exercise demonstrates the need of accurately characterizing the instrumental system before great photometric accuracies of Type Ia supernovae can be claimed. It also shows that this effect can have important astrophysical consequences since a small systematic shift of 0.02 mag in the $B-V$ color can introduce a $0.08$ mag error in the extinction corrected peak $B$ magnitudes of a supernova and thus lead to biased cosmological parameters. The data for the Type~Ib/c SN~1999ex present us with the first ever observed shock breakout of a supernova of this class. These observations show that shock breakout occurred 18 days before $B_{max}$ and support the idea that Type~Ib/c supernovae are due to core collapse of massive stars rather than thermonuclear disruption of white dwarfs. ",
        "introduction": "In recent years significant effort has gone into searching for and observing Type~Ia supernovae (hereafter referred to as SNe) with the purpose of determining cosmological parameters. Despite this progress, most of SNe~Ia observations have been limited to optical wavelengths, and relatively little is still known about the infrared (IR) properties of these objects. IR photometry of SNe has been limited to  a handful of events, including the early work of \\citet{elias81,elias85} and \\citet{frogel87}, and the most recent efforts by \\citet{jha99}, \\citet{hernandez00}, and \\citet{krisciunas00}. Combined IR studies of SNe~Ia \\citep{meikle00} have shown that SNe~Ia display a scatter in peak absolute magnitude no larger than $\\pm0.15$ mag in all three IR bands $JHK$, which is comparable to that obtained in the optical. To further exploit the usefulness of these objects as distance indicators, as well as to better understand the explosion mechanism and nature of SNe, more IR observations are clearly required.  Today the advent of new IR detectors enables us to obtain high quality observations and enlarge the hitherto small samples of SNe observed at these wavelengths. The ``Supernova Optical and Infrared Survey'' (SOIRS) was initiated in 1999 to make extensive IR and optical observations of bright candidates from the El Roble SN survey \\citep{maza81}, and from the Nearby Galaxies Supernova Search \\citep{strolger99}. Once objects were identified to be SN candidates, followup observations were conducted at several observatories including the Cerro Tololo Inter-American Observatory (CTIO) for optical/IR photometry, the Las Campanas Observatory (LCO) for IR photometry, and the European Southern Observatory (ESO) at Cerro Paranal and La Silla for optical/IR spectroscopy. During 1999-2000 this campaign resulted in a detailed study of 18 SNe. Of these events 10 were SNe~Ia, 7 SNe~II and 1 a Type Ib/c event.  This paper presents optical photometry for the two best-observed  SNe in this survey, SN~1999ee and SN~1999ex. SN~1999ee was discovered by M. Wischnjewsky on a T-Max 400 film obtained by L. Gonz\\'alez on 1999 October 7.15 UT (JD 2,451,458.65) in the course of the El Roble survey \\citep{maza99}. This SN exploded in IC~5179, a very active star-forming Sc galaxy with a heliocentric redshift of 3,498 $km$ $s^{-1}$. Adopting a value for $H_{\\circ}$=63.3$\\pm$3.5 $km$ $s^{-1}$ $Mpc^{-1}$ \\citep{phillips99} along with a redshift of 3,239$\\pm$300 $km$ $s^{-1}$ in the Cosmic Microwave Background (CMB) Reference System, IC~5179 is located at a distance of 51.2$\\pm$5.5 Mpc ($\\mu$=33.55$\\pm$0.23), which we adopt throughout this paper. An optical spectrum taken on 1999 October 9.10 revealed that SN~1999ee was a very young SN~Ia, which led us to give it a high priority among our list of targets for followup observations.  In a rare occurrence IC~5179 produced a second SN within a few weeks from the discovery of SN~1999ee. The second SN (1999ex) was discovered on 1999 November 9.51 UT by \\cite{martin99} at Perth Observatory through the course of the PARG Automated Supernova Search. Initially we classified it as a Ic event \\citep{hamuy99} based on the close resemblance to SN~1994I \\citep{filippenko95}, although we now believe it belongs to an intermediate Ib/c class \\citep{hamuy02}. Although SN~1999ex was discovered on 1999 November 9, examination of our data revealed that the object was present on our CCD images obtained 10 days before discovery. Further examination of the data showed that we had unintentionally detected the initial shock breakout resulting from the core bounce of SN~1999ex which, until now, has never been observed in a Ib/c event.  The observations gathered for SNe~1999ee and 1999ex have an unprecedented temporal and  wavelength coverage and afford the opportunity to carry out a detailed comparison with explosion and atmosphere models. In this paper we report final reduced $UBVRIz$ lightcurves for both of these SNe. The photometry presented here is available in electronic form to other researchers.\\footnote{At http://www.ociw.edu/$\\sim$mhamuy/lightcurves.html.} Section \\ref{obs.sec} discusses the observations and the data reduction process and Section \\ref{results.sec} presents final results for both events. We proceed with a discussion in Section \\ref{disc.sec} after which we summarize our conclusions.  Readers are to refer to K. Krisciunas et al. (2002, in preparation) and \\citet{hamuy02}, which include the analysis of IR photometry and optical/IR spectra, respectively. ",
        "conclusions": "\\label{conc.sec}  In a rare occurrence two SNe were observed simultaneously in IC~5179 at a redshift of z = 0.01 between 1999 October-December. For each image of these SNe we performed galaxy subtraction of late time images of their host galaxy after which we extracted $UBVRIz$ magnitudes. We obtained well-sampled lightcurves including the first $z$ band observations for a SN~Ia and SN~Ib/c.  Observations of the Type Ia SN~1999ee span from 10 days prior to $B_{max}$ to 53 days afterwards. SN~1999ee was characterized by a slow decline rate of $\\Delta$$m_{15}(B)$=0.94$\\pm$0.06. We estimate a value of host-galaxy reddening of 0.28$\\pm$0.04 which we use to derive reddening-free peak absolute magnitudes of $M_{B}$=-19.85$\\pm$0.28, $M_{V}$=-19.87$\\pm$0.26, and $M_{I}$=-19.43$\\pm$0.24. SN~1999ee is among the most luminous SNe~Ia and fits well in the peak luminosity-decline rate relation for SNe~Ia.  Near maximum we find systematic differences $\\sim$0.05 mag in photometry measured with two different telescopes, even though the photometry is reduced to the same local standards around the supernova using the specific color terms for each instrumental system. We use models for our bandpasses and spectrophotometry of SN~1999ee to derive magnitude corrections (S-corrections) and remedy this problem. This exercise demonstrates the need of accurately characterizing the instrumental system before great photometric accuracies of Type Ia supernovae can be claimed. It also shows that this effect can  have important astrophysical consequences since a small systematic shift of 0.02 mag in the $B-V$ color can   introduce a $0.08$ mag error in the extinction corrected peak $B$ magnitudes of a supernova and thus lead to biased cosmological parameters.  The data for the second SN observed in this galaxy -- Type~Ib/c SN~1999ex -- present us with the first ever observed shock breakout of a supernova of this class. These observations show that shock breakout occurred 18 days before $B_{max}$ and support the idea that Type~Ib/c supernovae are due to core collapse of massive stars rather than thermonuclear disruption of white dwarfs. The comparison of the lightcurves of SN~1999ex to other SNe~Ic events like SN~1983V and SN~1994I reveals large photometric differences among this class of objects, probably due to variations in the properties of the envelopes of their progenitors and/or explosion energies. Future theoretical modeling of this event along with spectral analysis and construction of a bolometric lightcurve will provide insight on relevant parameters describing its progenitor."
    },
    "0206/astro-ph0206324_arXiv.txt": {
        "abstract": "The large, powerful Fanaroff-Riley class II radio source 3CR 68.1 is an optically red quasar with strong evidence for dust reddening, intrinsic ultraviolet line absorption, and X-ray absorption.  In the context of its large, extended radio structure and the evidence for intrinsic material along the line of sight, it is a good candidate for a quasar seen through the edge of an obscuring torus proposed by unified schemes.  The compact radio core coincident with the optical position of the quasar is extremely weak compared to the very bright and luminous radio lobes, suggesting that the radio core might also suffer absorption.  We observed 3CR 68.1 at six frequencies and show that the radio core has a flat spectrum with no evidence for free-free absorption or other strong absorption mechanisms.  This result helps establish an empirical intrinsic lower range for core-to-lobe radio flux in the most powerful radio sources. ",
        "introduction": "The extreme properties of 3CR~68.1 indicate it is seen in an edge-on geometry (Brotherton et al. 1998), along a line of sight through dusty, ionized gas perhaps part of an obscuring torus (e.g., Antonucci 1993). Such a structure seen in such a geometry helps to explain the dust-reddened optical continuum (A$_V$ = 1.7 mag, Simpson \\& Rawlings 2000), scattered light leading to a rising continuum polarization up to 13\\% at 2000 \\AA\\ rest-frame (Brotherton et al. 1998), intrinsic ultraviolet absorption lines including metastable He I* (Brotherton et al. 1998), extremely weak and probably absorbed X-ray emission (Bregman et al. 1985), and the large radio lobe separation (53\\arcsec\\ at z=1.228) and small radio core dominance (Bridle et al. 1994) that mark an edge-on radio structure (e.g. Orr \\& Browne 1982).  The radio core of 3CR 68.1 is an order of magnitude fainter than that of any other 3CR quasar (1.1 mJy at 5 GHz, Bridle et al. 1984) and it is weak even among Fanaroff-Riley type II (FR II; Fanaroff \\& Riley 1974) radio galaxies. At the same time, the radio lobes make 3CR 68.1 one of the three most powerful 3CR quasars at 178 MHz in the compilation by Laing et al. (1983). We hypothesized that the radio core might be so faint because it suffers free-free absorption from the obscuring material along the line of sight, measurements of which would provide an additional probe of its physical conditions.  We obtained multi-band observations of 3CR~68.1 using the NRAO\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} Very Large Array (VLA) in order to test this hypothesis.  Following sections of this paper discuss the observations and data reduction (\\S 2), measurements of the core (\\S 3), our null result and its implications (\\S 4).  The final section (\\S 5) summarizes our conclusions. ",
        "conclusions": "While 3CR~68.1 clearly has dust and gas along the line of sight that causes absorption at optical and X-ray wavelengths, there does not appear to be any significant absorption toward the radio core at GHz frequencies.  The radio core spectrum is consistent with a flat ($\\alpha = 0$) spectrum of approximately 1 mJy, less than 1/1000 as bright as the radio lobes at 5 GHz, rest-frame, with no evidence for free-free or other absorption at radio frequencies.  The case of 3CR~68.1 apparently represents the intrinsic low extreme of the radio core dominance parameter used to characterize orientation in unified schemes."
    },
    "0206/astro-ph0206168_arXiv.txt": {
        "abstract": "{The K20 survey is an ESO VLT optical and near-infrared spectroscopic survey aimed at obtaining spectral information and redshifts of a complete sample of about 550 objects to $K_s\\leq20.0$ over two independent fields with a total area of 52 arcmin$^2$. In this paper we discuss the scientific motivation of such a survey, we describe the photometric and spectroscopic properties of the sample, and we release the $K_s$-band photometric catalog. Extensive simulations showed that the sample is photometrically highly complete to $K_s=20$. The observed galaxy counts and the $R-K_s$ color distribution are consistent with literature results. We observed spectroscopically 94\\% of the sample, reaching a spectroscopic redshift identification completeness of 92\\% to $K_s\\leq20.0$ for the observed targets, and of 87\\% for the whole sample (i.e. counting also the unobserved targets). Deep spectroscopy was complemented with multi-band deep imaging in order to derive tested and reliable photometric redshifts for the galaxies lacking spectroscopic redshifts. The results show a very good agreement between the spectroscopic and the photometric redshifts with $<z_{spe}-z_{phot}>=0.01$ and with a dispersion of $\\sigma_{\\Delta z}$=0.09. Using both the spectroscopic and the photometric redshifts, we reached an overall redshift completeness of about 98\\%. The size of the sample, the redshift completeness, the availability of high quality photometric redshifts and multicolor spectral energy distributions make the K20 survey database one of the most complete samples available to date for constraining the currently competing scenarios of galaxy formation and for a variety of other galaxy evolution studies. ",
        "introduction": "Contrary to surveys for high-$z$ galaxies selected in the optical, which are more sensitive to the star-formation activity, the selection of galaxies in the $K$-band has the important advantage of not being affected by strong K-correction effects. This is due to the similarity of the different galaxy spectral types in the near-infrared. Cowie et al. (1994) showed that over a broad range of redshifts the K-correction amplitude in the $K$-band is much smaller than those in $I$- or in $B$- bands, and very little dependent on the galaxy spectral types (from elliptical to irregular galaxies). Thus, the great advantage of a $K$-band selection is that the resulting samples are almost free from strong selection effects and do not critically depend on the galaxy types as in optical samples.  However, deriving the complete redshift distribution of a sample of faint $K$-band selected galaxies is a challenging goal even for 8-10m class telescopes. Indeed, a fraction of the galaxies are beyond the spectroscopic limits (e.g. very red galaxies with $R>26$) and/or lie in a redshift range where the spectra do not present prominent features in the observer frame (e.g. the redshift ``desert'' at $1.4<z<2.0$ for optical spectroscopy).  Songaila et al. (1994) carried out an optical spectroscopic survey of $K$-selected galaxies at different magnitude depths over a wide range of field sizes (from $\\sim$0.4 deg$^2$ down to $K$ \\ltsima 15, to $\\sim$5 arcmin$^2$ to $K$ \\ltsima 20) reaching an identification completeness of $\\sim$70\\% at $K\\sim 19-20$. Cowie et al. (1996) performed with the Keck I 10m telescope a deeper optical spectroscopic survey of galaxies with $K<20$ over a field of 26.2 arcmin$^2$, reaching a spectroscopic identification completeness of $\\sim$60\\% at $K\\sim 20$. Cohen et al. (1999a, 1999b) observed a sample of 195 objects over an area of 14.6 arcmin$^2$ reaching a spectroscopic completeness of 83.6\\% to $K_s<20$. A large spectroscopic survey is currently in progress over four independent fields covering a total area of about 100 arcmin$^2$ to $K<20$ (Stern et al. 2001).  In addition to purely spectroscopic surveys, a number of more recent projects made use mostly of photometric redshifts. For instance, Drory et al. (2001) selected $\\sim$5000 galaxies to $K<19$ over 998 arcmin$^2$, and estimated the photometric redshifts for 94\\% of them using $VRIJK$ photometry (spectroscopic redshifts were available for only $\\sim$6\\% of the sample). Other surveys based largely on photometric redshifts range from that of Firth et al. (2002), who covered a field of 744 arcmin$^2$ to $H<20.0-20.5$ and derived photometric redshifts for about 4000 galaxies using $UBVRIH$ photometry, to those of e.g. Fontana et al. (1999,2000) and Rudnick et al. (2001) based on very deep observations (typically to $K_s<21-23$) of very small fields (a few arcmin$^2$ ).  Surveys for faint galaxies selected in the $K$-band are very important to investigate the formation and evolution of galaxies.  One of the main questions of galaxy evolution is whether massive galaxies (e.g. ${\\cal M}_{stars}$\\gtsima $10^{11}$ M$_{\\odot}$) are the late product of merging of pre-existing disk galaxies occurring mostly at $z<1.5-2$, as predicted by the current CDM scenarios (e.g. Kauffmann et al. 1993, 1996; Cole et al. 2000; Baugh et al. 1996; Somerville et al. 2001, Baugh et al. 2002, and references therein), or whether they formed at $z>2-3$ during a short-lived and single event of vigorous star formation, followed by a passive evolution (or pure luminosity evolution, PLE) of the stellar population to nowadays (see e.g. Renzini 1999; Renzini \\& Cimatti 1999; Peebles 2002, for recent reviews).  Since the near-IR light is a good tracer of a galaxy stellar mass (Gavazzi et al. 1996; Madau, Pozzetti \\& Dickinson 1998; Kauffmann \\& Charlot 1998, KC98 hereafter), the above scenarios can in principle be tested by selecting galaxies at 2.2$\\mu$m in the $K$-band and deriving their redshift distribution and physical and evolutionary properties (e.g. Broadhurst et al. 1992).  KC98 estimated that $\\sim 60\\%$ and $\\sim 10\\%$ of the galaxies in a $K<20$ sample are expected to be at $z>1$, respectively, in a PLE and in a standard CDM hierarchical merging model (cf. their Fig. 4). Such a large difference was in fact one of the main motivations of our original project to undertake a complete redshift survey for all objects with $K<20$ in a small area of the sky. However, more recent models consistently show that for $z>1$ the difference between the predictions of different scenarios is less extreme than in the KC98 realization (Menci et al., Pozzetti et al., in preparation; see also Firth et al. 2002). Part of the effect is due to the now favored $\\Lambda$CDM cosmology which pushes most of the merging activity in hierarchical models at earlier times compared to $\\tau$CDM and SCDM models, and therefore get closer to the PLE case. Moreover, a different tuning of the star-formation algorithms (to accomodate for more star formation at high-$z$) also reduces the differences between the two scenarios (e.g. Somerville et al. 2001).  In order to address the question of the formation and evolution of galaxies and to constrain the current models, we performed a new spectroscopic survey of $K$-selected galaxies with the ESO VLT. In this paper, we present the main scientific motivations, we define the sample selection, we discuss the photometry, the completeness, we present the main spectroscopic and photometric redshift results, and we release the $K_s$-band photometric catalog. The spectral and multicolor catalogs will be presented elsewhere together with the forthcoming papers on the scientific analysis of the sample. Magnitudes are given in Vega system. The widely accepted cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$ is adopted throughout this paper. ",
        "conclusions": "We discussed the general scientific background and aims of the K20 survey and we described the sample photometric and spectroscopic properties.  The relevant advantages of the K20 database are the small K-correction effects due to the $K_s$-band selection, its size (about 500 galaxies), the distribution of the targets in two independent fields, the use of near-IR spectroscopy for a subsample of the targets, and the availability of a large deep imaging database from the optical to the near-IR.  Extensive simulations show that the sample is highly complete to $K_s=20.0$ and not affected by strong selection effects on the galaxy population, with the possible exception of a slight underestimate of the total flux for large and luminous early-type galaxies with de Vaucouleurs surface brightness profiles. Such a selection effect can be taken into account when comparing the observed K20 sample with model predictions. The observed galaxy counts and the $R-K_s$ color distribution are in agreement with literature results.  Optical spectroscopy, aided by near-IR spectroscopy, allowed us to achieve a high redshift completeness for a sample of galaxies selected in the near-IR (94\\% to $K_s<19.0$ and 87\\% $K_s<20.0$), and to obtain spectroscopic redshifts in the ``desert'' of $1.4<z<2.2$.  The high-quality deep imaging database allowed us also to obtain tested and reliable photometric redshifts for the unobserved or spectroscopically unidentified galaxies. Using the photometric redshifts, the global completeness ($z_{spec}+z_{phot}$) increases to 98.3\\% of the total sample.  Such a rich photometric and spectroscopic database makes the K20 sample a key tool to investigate the formation and evolution of galaxies."
    },
    "0206/astro-ph0206442_arXiv.txt": {
        "abstract": "We have carried out additional spectroscopic observations in the field of cluster Cl 1324+3011 at $z = 0.76$. Combined with the spectroscopy presented in Postman, Lubin \\& Oke (2001, AJ, 122, 1125), we now have spectroscopically confirmed 47 cluster members. With this significant number of redshifts, we measure accurately the cluster velocity dispersion to be $1016^{+126}_{-93}$ km s$^{-1}$. The distribution of velocity offsets is consistent with a Gaussian, indicating no substantial velocity substructure. As previously noted for other optically-selected clusters at redshifts of $z \\gtrsim 0.5$, a comparison between the X-ray luminosity ($L_x$) and the velocity dispersion ($\\sigma$) of Cl 1324+3011 implies that this cluster is underluminous in X-rays by a factor of $\\sim 3 - 40$ when compared to the $L_x - \\sigma$ relation for local and moderate-redshift clusters. We also examine the morphologies of those cluster members which have available high-angular-resolution imaging with the {\\it Hubble Space Telescope} (HST).  There are 22 spectroscopically-confirmed cluster members within the HST field-of-view.  Twelve of these are visually classified as early-type (elliptical or S0) galaxies, implying an early-type fraction of $0.55^{+0.17}_{-0.14}$ in this cluster. This fraction is a factor of $\\sim 1.5$ lower than that observed in nearby rich clusters. Confirming previous cluster studies, the results for cluster Cl 1324+3011, combined with morphological studies of other massive clusters at redshifts of $0 \\le z \\lesssim 1$, suggest that the galaxy population in massive clusters is strongly evolving with redshift. This evolution implies that early-type galaxies are forming out of the excess of late-type (spiral, irregular, and peculiar) galaxies over the $\\sim 7$ Gyr timescale. ",
        "introduction": "Clusters of galaxies represent the largest scale of fully collapsed, virialized structures in the universe and provide a powerful probe of the evolution of structure formation via dynamical collapse. Therefore, quantifying the abundance and dynamical state of clusters is key to understanding the formation of clusters and the evolution of the galaxies within them. With the advent of sensitive X-ray and optical satellites, such as {\\it Einstein}, ROSAT, Chandra, XMM, and HST, and large ground-based telescopes, such as the Keck 10-m and the VLT 8-m, it has been possible to extend cluster studies up to redshifts of $z \\sim 1$.  Optical surveys have revealed a only modest decline in the comoving volume density of rich clusters out to $z \\sim 1$ (Gunn, Hoessel \\& Oke 1986; Postman et al.\\ 1996; Carlberg et al.\\ 1997; Bahcall \\& Fan 1998; Postman et al.\\ 2002).  X-ray observations indicate that the group/cluster X-ray luminosity function does not evolve significantly to $z \\sim 0.8$ for modest X-ray emitters in the luminosity range of $2 \\times 10^{42}$ to $3 \\times 10^{44}~h_{50}^{-2}~{\\rm ergs~s^{-1}}$ (Ebeling et al.\\ 1997; Nichol et al.\\ 1997; Rosati et al.\\ 1998; Jones et al.\\ 1998); however, there appears to be a measurable decrease in the number of very X-ray luminous ($L_{x} \\gtrsim 4 \\times 10^{44}~h_{50}^{-2}~{\\rm ergs~s^{-1}}$) clusters out to $z \\sim 0.9$ (Edge et al.\\ 1990; Henry et al.\\ 1992; Bower et al.\\ 1994; Jones et al.\\ 1998; Henry 2002; Lewis et al.\\ 2002). With these large-area surveys, it has been possible to study the global X-ray and optical properties of clusters and use the observed evolution, or lack thereof, to place strong constraints on the cosmological world model (e.g., Luppino \\& Gioia 1995; Carlberg et al.\\ 1997; Bahcall \\& Fan 1997; Borgani et al.\\ 2001).  The galaxy populations within massive clusters, and how their properties evolve with time, have also been the subject of considerable study.  In local clusters, early-type (elliptical and S0) galaxies comprise more than 80\\% of the total population.  The galaxy content of clusters is part of the general morphology--density relation which describes the strong correlation between galaxy morphology and environment. As the local density increases, the fraction of early-type galaxies increases, while the fraction of late-type galaxies decreases (Dressler et al.\\ 1980a,b; Postman \\& Geller 1984). Studies of clusters at $z < 0.6$ have revealed deviations in this universal relation because the morphology and color of the cluster members are evolving.  One of the most notable changes is the progressive bluing of the cluster's galaxy population with redshift. Butcher \\& Oemler (1984) found that the fraction of blue galaxies in a cluster is a strong function of redshift, increasing from $<5\\%$ in local clusters to $\\sim 20\\%$ in clusters at $z \\sim 0.5$.  HST imaging reveals that many of these blue galaxies are either normal spirals or have peculiar morphologies, resulting in late-type fractions which are $\\sim 3$ times higher than the average current epoch cluster (Dressler et al.\\ 1994, 1997; Couch et al.\\ 1994; Oemler, Dressler \\& Butcher 1997; Fabricant, Franx \\& van Dokkum 2000).  Even though these clusters show an increased fraction of blue galaxies, moderate-redshift clusters still contain a population of early-type galaxies distinguished by extremely red colors and a tight color--magnitude (CM) relation (a ``red envelope''). Both the mean color and the CM relation are consistent with that of present--day ellipticals (Butcher \\& Oemler 1984; Arag\\'on-Salamanca et al.\\ 1991; Stanford, Eisenhardt \\& Dickinson 1995). At $z \\gtrsim 0.4$, the red envelope moves bluewards with redshift (e.g.\\ Arag\\'on-Salamanca et al.\\ 1993; Lubin 1996; Stanford et al.\\ 1997; Stanford, Eisenhardt \\& Dickinson 1998). This color trend is consistent with passive evolution of an old stellar population formed largely at $z \\gtrsim 2$. The red envelope is reasonably narrow and uniform from cluster to cluster, implying that the galaxies formed in an episode of star formation lasting no longer than a few Gyrs (Bower, Lucey \\& Ellis 1992a,b; Ellis et al.\\ 1997; Stanford et al.\\ 1998; Bower, Kodama \\& Terlevich 1998; van Dokkum et al.\\ 2001).  To probe the properties of clusters and their galaxy populations at even earlier stages in their development, we have conducted an extensive observational program to study nine candidate clusters of galaxies at $z \\gtrsim 0.7$ (Oke, Postman \\& Lubin 1998; hereafter Paper I) with the goal of establishing an evolutionary reference sample of clusters analogous to the MORPHS survey at $z \\sim 0.5$ (Smail et al.\\ 1997). The cluster sample was chosen from the optical/near-IR surveys of Gunn, Hoessel \\& Oke (1986) and Postman et al.\\ (1996). For each cluster, we have obtained deep $BVRIK'$ photometry from the Keck 10-m and KPNO 4-m telescopes, over 120 low-resolution spectra from a $R$-band magnitude-limited spectroscopic survey with Keck, and high-angular-resolution imagery from the Wide Field Planetary Camera 2 (WFPC2) aboard HST. Five papers based on this survey have been published so far. They include Postman, Lubin \\& Oke (1998; hereafter, Paper II) and Lubin et al.\\ (1998; hereafter, Paper III) which describe the photometric, spectroscopic, and morphological properties of Cl 1604+4304 ($z = 0.90$) and Cl 0023+0423 ($z = 0.83$). In addition, Postman, Lubin \\& Oke (2001; hereafter, Paper IV) describe the cluster dynamical properties and the spectrophotometric properties of the galaxies in Cl 1324+3011 ($z = 0.76$) and Cl 1604+4321 ($z = 0.92$).  In this paper, we present additional spectroscopic observations of the massive cluster, Cl 1324+3011 at $z = 0.76$, taken at the Keck 10-m telescope. The goal of these observations is to obtain more cluster redshifts in order to measure the dynamics of the cluster and the properties of the cluster galaxy population more accurately. The new observations are described in \\S 2. We have used these new spectroscopic data, combined with those described in Paper IV, to improve the measurement of the cluster velocity dispersion and total mass (\\S 3), to make a more accurate comparion between the X-ray and optical properties of this cluster (\\S 4), and to estimate with reasonable uncertainties the morphological mix of the cluster galaxy population. The discussion of our results and their implication for cluster evolution is presented in \\S 6.  Unless otherwise noted, we adopt $h_{65} \\equiv H_o/(65~{\\rm km~s^{-1}~Mpc^{-1}}) = 1$, $\\Omega_o = 0.2$, and $\\Lambda = 0$. ",
        "conclusions": "We have obtained additional spectroscopy for galaxies in the field of the massive cluster Cl 1324+3011 at $z = 0.76$. With these new data and those previously taken as part of our original cluster survey (Paper IV), we have now spectroscopically-confirmed 47 cluster members, of which 22 have available high-angular-resolution HST imaging and, therefore, have measured morphological properties.  We have used these data to measure accurately the cluster velocity dispersion ($\\sigma$) and the fraction of early-type galaxies in the cluster galaxy population ($f_{E+S0}$).  We find $\\sigma = 1016^{+126}_{-93}$ km s$^{-1}$ and $f_{E+S0} = 0.55^{+0.17}_{-0.14}$. Both of these quantities imply significant differences between Cl 1324+3011 and the average properties of local clusters of galaxies.  Firstly, Cl 1324+3011 does not follow the relation between X-ray luminosity and velocity dispersion observed in nearby and moderate-redshift clusters. We find that the X-ray luminosity of Cl 1324+3011 is low by a factor of $\\sim 3 - 40$ for its velocity dispersion and, therefore, its estimated total mass. This cluster is typical of optically-selected clusters at redshifts of $z \\gtrsim 0.5$ and seems to imply that the galaxies and the gas are no longer in thermal equilibrium. We expect such behavior if these clusters are still dynamically young. For example, at these early epochs, clusters of galaxies may still be in the process of merging.  This process can result in a velocity dispersion which is inflated due to infalling matter or large scale structure, artificially implying a high mass. However, the X-ray emission can also evolve in a significant way if (1) subclumps in the process of cluster formation are not yet fully virialized or (2) subclumps are currently in the process of merging. In the former case, the gas will reside in lower-temperature (and lower-density) subclumps, substantially reducing its X-ray emission relative to the total mass within the region. In the latter case, the cluster would have a significant excess luminosity and temperature over that of non-merging clusters, which results from higher-density shocks between the merging cores (Ricker 1998; Ricker \\& Sarazin 2001; Sarazin 2001). These enhancements can be over 4--10 times in lumonisity and up to 3 times in temperature, and the effect can last at least one crossing time ($\\sim$ 1--2 Gyr).  There is, in fact, strong observational evidence that a large fraction (up to 75\\%) of clusters at $z \\sim 1$ are still in the process of forming and are, therefore, dynamically young (see e.g.\\ Henry 2002). The signatures of this formation include double peaks, strong substructure, and/or a filamentary appearance in the distribution of galaxies, gas, and total mass (e.g., Lubin \\& Postman 1996; Henry et al.\\ 1997; Luppino \\& Kaiser 1997; Donahue et al.\\ 1998; Gioia et al.\\ 1999; Della Ceca et al.\\ 2000; Ebeling et al.\\ 2000; Jeltema et al.\\ 2001; Stanford et al.\\ 2001; Hashimoto et al.\\ 2002).  Despite evidence suggesting that most clusters at $z > 0.7$ are not relaxed, there is no evidence that the average relation between X-ray luminosity and temperature has evolved with redshift. Using a sample of 11 X-ray--selected clusters, Holden et al.\\ (2002) find no statistically significant evolution in the slope or zero-point of the $L_x-T_x$ relation at a median redshift of $z = 0.83$ (see also Donahue et al.\\ 1999; Della Ceca et al.\\ 2000; Fairley et al.\\ 2000; Arnaud, Aghanim \\& Neumann 2002).  However, there may be preliminary indications that some clusters (for example, the double-peaked cluster RXJ 0152.7-1357 at $z = 0.833$; Ebeling et al.\\ 2000; Della Ceca et al.\\ 2000; Lewis et al.\\ 2002) do fall well off the $L_x-T_x$ relation.  This behavior would be expected if the dynamical and evolutionary state of clusters at these early epochs is diverse, as indicated by some optical, X-ray, and weak-lensing studies.  As of yet, we have not been able to determine where optically-selected clusters at high redshift fall on the $L_x-T_x$ relation as none have been the subject of X-ray spectral observations. However, this situation will soon be rectified for Cl 1324+3011 as we are currently in the process of analyzing a 40 ksec XMM observation of this cluster. For the first time, we will have accurate information on the luminosity, temperature, and mass of an optically-selected cluster at $z \\sim 0.8$. These data will go far in determining what physical processes are responsible for the differences between X-ray and optically-selected clusters at these redshifts.  Secondly, we measure an early-type fraction in Cl 1324+3011 which is lower by a factor of $\\sim 1.5$ compared to that observed in nearby rich clusters (Dressler 1980a,b; Andreon et al.\\ 1997). A decrease in the number of early-type galaxies was first observed by Dressler et al.\\ (1997), Andreon et al.\\ (1997), Fabricant et al.\\ (2000), and van Dokkum et al.\\ (2000) in massive clusters with redshifts between $z \\approx 0.3 - 0.8$. The results for Cl 1324+3011, combined with these previous studies, show a general trend of decreasing early-type fraction with redshift. The most significant question raised by this evolution is what physical mechanisms associated with the cluster environment can transform disk galaxies into spheroids. Over the years, at least three plausible mechanisms have been suggested. The first is galaxy harassment which involves a high-speed encounter of a faint spiral (Sc or Sd) galaxy with the cluster tidal field or a brighter cluster galaxy. The tidal shock strips the stellar disk, producing a remnant consistent with the mass profiles of dwarf ellipticals (Moore, Lake \\& Katz 1998).  The second is ram pressure stripping by the intracluster medium as an early spiral (Sa or Sb) galaxy enters the cluster potential. This interaction can result in either a burst or truncation of star formation (Gunn \\& Gott 1972; Abadi, Moore \\& Bower 1999; Poggianti et al.\\ 1999).  In both scenarios, the gas is eliminated, and the disk fades. The final mechanism is galaxy--galaxy merging. The morphological signatures of a merger between two disk galaxies, such a long, bright tidal tail, would be visible for only 1--2 Gyrs.  After this time, the tidal debris becomes indistinguishable from the main merger remnant, a morphologically--normal elliptical galaxy (Schweizer 1986; Hibbard et al.\\ 1994; Mihos 1995). Merging is unlikely to occur in the cluster environment because the velocities are too high; however, small groups of galaxies being accreted into the cluster are ideal sites for merging because the group velocity dispersion is on the order of the internal velocity of the member galaxies (Aarseth \\& Fall 1980; Barnes 1985; Merritt 1985).  At least two of these processes require infall of individual galaxies and/or galaxy groups into the cluster environment. Since the field population is comprised mainly of spiral and irregular galaxies (e.g\\ Driver et al.\\ 1995; Abraham et al.\\ 1996), infall is the natural means for providing a reservoir of late-type galaxies. Simulations predict that clusters should form at the intersections of large scale sheets and filaments (Shandarin \\& Zel'dovich 1989) through the accretion of smaller subclumps from the surrounding environment (Evrard 1990; Lacey \\& Cole 1993, 1994). Semi-analytic models show that the accretion rates increase strongly with look-back time. In early stages of cluster formation, the infall rate is up to a factor of $\\sim 4$ higher than in present-day clusters (Kauffmann 1995). Observationally, the strong substructure and filamentary nature in a large fraction of clusters at $z > 0.7$ (see above) suggests that clusters are actively forming at these redshifts.  These results indicate that infall regions are likely to be active sites of galaxy evolution.  Unfortunately, the field-of-view of WFPC2, which is only $160'' \\times 160''$, has strongly limited morphological studies in these regions. Because of these constraints, only one morphological study of a high-redshift cluster has extended beyond the central $0.5~h_{65}^{-1}$ Mpc (i.e.\\ a single WFPC2 pointing), that of the very X-ray luminous cluster MS 1054.4-0321 at $z = 0.83$ (van Dokkum et al.\\ 1999, 2000). This WFPC2 study, which covered a 2.7 $h_{65}^{-1} \\times$ 1.8 $h_{65}^{-1}$ Mpc$^{2}$ area in the F606W and F814W filters, has revealed some stunning results. The authors find that 17\\% of the $L \\ge L_{*}$ cluster population is comprised of ongoing mergers which will likely evolve into luminous ($\\sim 2 L_{*}$) early-type galaxies. These mergers are preferentially found in the outskirts of the cluster, indicating they probably reside in cold infalling clumps. If the galaxy population in MS1054.4-0321 is typical for its redshift, 50\\% of the early-type galaxies in present-day clusters have experienced a major merger at $z < 1$. The crucial question is whether MS1054.4-0321 is typical of other clusters at these redshifts.  In order to convincingly discriminate between the physical processes which can transform galaxy morphologies in this environment, it is necessary to study galaxy properties at large radii for a statistical sample of high-redshift clusters. We expect significant progress to be made on this front with the wide-area capability of the Advanced Camera for Surveys which was placed aboard HST this year."
    },
    "0206/astro-ph0206397_arXiv.txt": {
        "abstract": "{ We present our observations of the galaxy UGS 5600 with a long-slit spectrograph (UAGS) and a multi-pupil field spectrograph (MPFS) using  the 6-m telescope of  Special Astrophysical Observatory. Radial-velocity fields of the stellar and gaseous components were constructed for the central region and inner ring of the galaxy. We proved the existence of two almost orthogonal kinematic subsystems and conclude that UGC 5600 is a galaxy with an inner polar ring. In the circumnuclear region, we detected noncircular stellar motions and suspected the existence of a minibar. The emission lines are shown to originate in H II regions. We estimated the metallicity from the intensity ratio of the [NII]$\\lambda6583$ and \\Ha lines to be nearly solar, which rules out the possibility that the polar ring was produced by the accretion of gas from a dwarf companion. c 2002 MAIK \"Nauka/Interperiodica \". } ",
        "introduction": "Polar-ring galaxies (PRGs)  are objects with two kinematic subsystems rotating in roughly orthogonal planes.  Based on its structural features (see Fig. 1), Whitmore et al. (1990) listed the peculiar galaxy UGC 5600 among the most probable PRG candidates:  its amorphous main body is surrounded by a broad outer ring and the brightenings attributed to the inner ring are observed in the EW direction on both sides of the center at a distance of $\\sim10''$. UGC 5600 is a member of a double system (VV 330).   Its companion, the galaxy UGC 5609, is at 1\\farcm4 (about 15 kpc in projection onto the plane of the sky)  to the southeast and has a similar radial velocity.  Recently two other galaxies with similar redshifts have been detected near VV 330 (Galletta et al. 1997) ;  all of them may represent a group of galaxies.  Among other galaxies from the catalog of Whitmore et al. (1990)  , UGC 5600 was observed in the radio range at a wavelength of 21 cm (Richter et al.  1994).  The $M_{HI} /L{B}$ ratio was found to be 0.86 , characteristic of late-type galaxies. CO-line observations revealed molecular hydrogen in all structures of the galaxy (Galletta et al. 1997).  The spectra of the galaxy along its major and minor axes were obtained by Reshetnikov and Combes (1994).  The radial-velocity curves are complex in shape. The authors suggested counterrotation in the central part of the galaxy ($ r<5''$ )  and gas rotation around its major axis. In addition, they pointed out that the galaxy is rich in gas and that the \\Ha  emission extends to $30''$ (6 kpc).  A detailed photometric study of UGC 5600 (Karataeva et al., 2001) shows that this is most likely a late-type spiral (Scd)  galaxy with an inner polar ring, which is projected onto the galaxy main body on the northern side and is seen through it on the southern side. The structure that was taken in the catalog of Whitmore et al. (1990)  as an outer ring represents two tightly wound spiral arms.  The final conclusion whether UGC 5600 belongs to PRGs can be reached only after proving the existence of two nearly orthogonal kinematic subsystems.  The distance to the galaxy is 37.6 Mpc ($H_0 = 75 \\km \\mbox{Mpc}~{-1} $)  and the scale is 0.18 kpc in $1''$ . ",
        "conclusions": "In conclusion, we summarize our main results. \\begin{enumerate}  \\item Based on 2D spectroscopy, we constructed the radial velocity fields of the stellar and gaseous components for the central regions of the peculiar galaxy UGC 5600.  \\item An analysis of these fields revealed two kinematic subsystems: the first is related to the galactic disk and the second is related to the inner ring.  \\item The angle between the disk and ring planes was found about $78^\\circ\\pm5^\\circ$; i.e., the inner ring is polar. This provides compelling evidence that the galaxy UGC 5600 belongs to PRGs. \\item We established from the intensity ratio of forbidden and permitted lines that the emission originates in H II regions; the metallicity was estimated. It proved to be nearly solar, which rules out a dwarf galaxy as the donor in forming the polar ring. \\end{enumerate}"
    },
    "0206/astro-ph0206504_arXiv.txt": {
        "abstract": "We present images of the J=10--9 rotational lines of HC$_3$N in the vibrationally excited levels 1v7, 1v6 and 1v5 of the hot core (HC) in Orion KL. The images show that the spatial distribution and the size emission from the 1v7 and 1v5 levels are different. While the J=10--9 1v7 line has a size of $4''\\times 6''$ and peaks $1.1''$ NE of the 3 mm continuum peak, the J=10--9 1v5 line emission is unresolved ($<3''$) and peaks $1.3''$ south of the 3 mm peak. This is a clear indication that the HC is composed of condensations with very different temperatures (170 K for the 1v7 peak and $>230$ K for the 1v5 peak). The temperature derived from the 1v7 and 1v5 lines increases with the projected distance to the suspected main heating source I. Projection effects along the line of sight could explain the temperature gradient as produced by source I. However, the large luminosity required for source I, $>5\\, 10^{5}$ L$_\\odot$, to explain the 1v5 line suggests that external heating by this source may not dominate the heating of the HC. Simple model calculations of the vibrationally excited emission indicate that the HC can be internally heated by a source with a luminosity of $10^5$ L$_\\odot$, located $1.2''$ SW of the 1v5 line peak ($1.8''$ south of source I). We also report the first detection of high-velocity gas from vibrationally excited HC$_3$N emission. Based on excitation arguments we conclude that the main heating source is also driving the molecular outflow. We speculate that all the data presented in this letter and the IR images are consistent with a young massive protostar embedded in an edge-on disk. ",
        "introduction": "Located at only 450 pc from the sun, the Orion KL region has been the prototype source to study the formation of high mass stars. This region shows all the signatures of massive star formation. It contains a plethora of IR cores ($\\simeq 20$) \\citep{dougados93,gezari98}, H$_2$O masers \\citep{genzel81,gaume98}, millimeter continuum emission \\citep{plambeck95,blake96}, compact radio sources \\citep{churchwell87,menten95}, and molecular hot cores and fast molecular outflows \\citep{blake96,wright96,chandler97,rodriguez99b}.  It is believed that most of the IR sources in the Orion KL nebula are not self luminous but show reprocessed emission escaping through inhomegeneities of the dense material surrounding a deeply embedded source. Recently, \\citet{menten95} and \\citet{gezari98} proposed that the main heating source in the region is radio source I. The lack of an IR counterpart to source I has led some authors to invoke large intrinsic foreground extinction towards source I. However, \\citet{blake96},\\citet{wright96} and \\cite{chandler97} have shown that the dust and molecular peak emission, coincide neither with radio source I nor with IRc2. The main peak of continuum emission at 1.3 mm, is $1''$ east from source I. Furthermore, based on the distribution of the HC$_3$N J=24--23 line in the 1v7 vibrationally excited state, \\citet{blake96} have proposed that source I is the dominant energy source in the region, and that there is no evidence of internal heating within the molecular hot core (HC). However, \\citet{kaufman98} based on a model for the heating of the HC concluded that it is internally heated by young embedded stars.  In this letter we investigate the thermal structure of the Orion KL HC by using the J=10--9 rotational line of HC$_3$N in the vibrational excited levels 1v7, 1v6 and 1v5. If the hot core material is not internally heated and source I is the main heating source one would expect the emission from the HC$_3$N rotational lines in the vibrational states (hereafter HC$_3$N*) with larger excitation energy to peak close to source I. We find that the emission of the HC$_3$N J=10--9 lines in the different vibrational levels peak at different positions indicating that the HC is internally heated. ",
        "conclusions": "The origin of the heating of the HC and of the driving source of the molecular outflow in Orion KL has been the subject of debate during the last 10 years and it is not settled yet. IRc2 and later source I have been the most likely candidates. \\citet{kaufman98} have studied the heating of the HC by source I and conclude that the HC is internally heated by young stars. The detection of the largest excitation temperature derived from the 1v5/1v7 ratio, located $\\sim 2''$ south of source I, also suggests that this source might not be the main heating source. One could consider that the excitation temperature gradient can be due to projection effects and that HC1 is closer to source I than HC2. In this case, the minimum separation between the 1v5/1v7 maximum and source I would be the projected distance; ie $\\simeq$ 1000 AU. In order to externally heat the H$_2$ column densities of HC1 of $\\sim 2\\, 10^{23} {\\rm cm^{-2}}$ up to temperatures of 230 K, source I should have a luminosity of $> 5\\, 10^{5} L_\\sun$ \\citep{kaufman98}. This is larger than the total luminosity of the KL region \\citep{werner76}. Heating of HC1 by the combination of sources I and n will not substantially change the required total luminosity of source I since source n is considered to be less luminous than source I and is located at the same projected distance from HC1 as source I. IR radiation from the shocks generated in the molecular outflows can be ruled out since the mechanical luminosity of the outflow is only $10^2\\ L_\\sun$ \\citep{rodriguez99}. We conclude that the most likely explation for the heating of HC1 is internal heating by young embedded stars. The lack of radio emission from HC2 might be an indication of mass accretion \\citep{walmsley95}.   The high luminosity ($\\sim 10^5$ L$_\\odot$) IR source heating HC1 is, like source I, a very weak emitter at 20$\\mu$m. The 1v5 line emission is anticorrelated with the IR emission at 20$\\mu$m measured by \\citet{gezari98} with $1''$ resolution. The morphology of the IR emission is possibly influenced by foreground extinction \\citep{wright96}, and HC1 can be obscured by the presence of foreground cold material. From the parameters in Table 1 for T$_{max}$, the expected flux density within $1''$ at 20$\\mu$m from HC1 should be $4\\, 10^{3}$ Jy, a factor of $\\approx 500$ larger than the observed flux. To match the observed and predicted fluxes within $1''$, the IR source in HC1 must be hidden behind 6 magnitudes of extinction at 20 $\\mu$m, i.e. an H$_2$ column density of $1.2\\, 10^{23}\\, {\\rm cm^{-2}}$. As expected, the cold gas in front of the IR source not observed in the HC$_3$N* lines is 20\\% of the HC1 total column density of $4\\, 10^{23}\\, {\\rm cm^{-2}}$, derived from the dust continuum emission \\citep{blake96} also measured with $1''$ resolution. Then, the IR map at 20 $\\mu$m seems to be consistent with our estimates for the luminosity of the IR source embedded in HC1.  Source I is the driving source of the SiO outflow, but is located just outside the HC. Furthermore, the molecular outflow observed in HC$_3$N*  seems to be driven by the IR source in HC1. These facts suggest, as proposed by \\citet{rodriguez99}, that the Orion KL region contains multiple molecular outflows.  The data presented in this paper, the presence of a high luminosity IR source embedded in HC1 powering the high-velocity gas observed in HC$_3$N* emission and the properties of the IR emission in the Orion KL region (morphology, luminosity and polarization) can be explained in the scenario of a young massive star deepely embbeded in molecular material \\citep{plambeck82}. The circumstellar material must be very anisotropic since it is located in the narrow ridge of minimum IR emission with strong molecular line emission and has large column densities along the line of sight. Although our angular resolution is not enough to resolve the circumstellar material, the most likely morphology is that of an edge-on disk. The circumstellar material is internally heated by young massive protostars with a luminosity of $10^5$ L$_\\odot$. The IR radiation heats the most inner parts giving rise to the emission observed in HC$_3$N*. The IR radiation escapes in the direction of the lowest column densities (perperdicular to the ridge) illuminating most of the IR sources as shown by the polarization vectors which point towards a location in the vicinity of HC1 \\citep{werner83}. The massive protostar drives the large molecular outflow observed in CO that is collimated by the circumstellar material. The material very close to the IR source where the mass ejection originates is observed as high velocity wings in HC$_3$N*.  This work has been partially supported by the Spanish Ministerio de Ciencia y Tecnologia grant ESP-4519-PE and SEPCT grant AYA2000-927. We want to thank the anonymous referees for their constructive comments which helped to improve substantially the original version of this letter."
    },
    "0206/astro-ph0206218_arXiv.txt": {
        "abstract": "Accretion onto a compact object must occur through a disc when the material has some initial angular momentum. Thin discs and the thicker {\\em low radiative efficiency accretion flows} (LRAFs) are solutions to this problem that have been widely studied and applied. This is an introduction to these accretion flows within the context of X-ray binaries and cataclysmic variables. ",
        "introduction": "\\subsection{Accretion efficiency, Eddington luminosity} In Newtonian dynamics, the available gravitational energy from a small mass $m$ moved from infinity to the surface $R_\\star$ of the accretor is $G M_\\star m / R_\\star$. The energy released can be a sizeable fraction $\\eta$ of the rest-mass energy $m c^2$ when the accretor is a neutron star or black hole: assuming a thin disc and free fall at the innermost stable circular orbit then $\\eta\\approx 0.42$ for a maximally rotating Kerr black hole (\\citealt{nt}; see also \\citealt{gammie}).  The Eddington luminosity \\ledd{} is the luminosity for which radiation pressure exactly balances the gravitational pull on the accreted material. Accretion is quenched above this luminosity \\citep[but see][]{begelman,shaviv}. For a spherical inflow of ionised hydrogen such that the radiation pressure is due to Thomson scattering on electrons and the dominant gravitational pull is on the protons, \\ledd{} is: \\begin{equation} \\frac{GM_\\star {\\rm m}_{\\rm p}}{R^2}= \\frac{{\\rm L}_{\\rm E}}{4\\pi R^2} \\frac{\\sigma_{\\rm T}}{c} {\\rm ~~hence~~} {\\rm L}_{\\rm E}\\approx 10^{38} ~(M_\\star/{\\rm M}_\\odot) {\\rm ~erg}\\cdot{\\rm s}^{-1} \\end{equation} The luminosities of XRBs and CVs are generally lower than this limit \\citep{xrbs,cvs}. The Eddington mass accretion rate $\\dot{\\rm M}_{\\rm E}$ is defined as: \\begin{equation} \\eta\\dot{\\rm M}_{\\rm E}c^2={\\rm L}_{\\rm E} {\\rm ~~i.e.~~} \\dot{\\rm M}_{\\rm E}\\approx 10^{18} ~(M_\\star/{\\rm M}_\\odot)(0.1/\\eta) {\\rm ~g~s}^{-1} \\end{equation}  \\subsection{Accretion flow temperature} Assuming the gravitational energy is released into radiation at the surface of the compact object, the minimum temperature of the flow is that of the black body radiating the same luminosity. At the Eddington limit this gives: \\begin{equation} 4\\pi R^2_\\star \\sigma T^4_{\\rm bb}={\\rm L}_{\\rm E} {\\rm ~~i.e.~~} {\\rm k}T_{\\rm bb}\\approx 1.5 ~{\\rm keV}~({\\rm L}_{\\rm E}/10^{38}{\\rm erg~s}^{-1})^{1/4} (R_\\star/ 10{\\rm ~km})^{-1/2} \\end{equation} Accreting black holes and neutron stars should radiate mainly in soft X-rays while white dwarfs ($R_\\star\\approx 10^4$~km) should radiate in UV. This is consistent with observations \\citep{xrbs,cvs}.  On the other hand, the maximum temperature is obtained when all the gravitational potential energy is transformed into thermal energy $e$ without radiation losses (adiabatic flow): \\begin{equation} e=\\frac{3}{2}\\frac{{\\rm k}T_{\\rm g}}{\\mu m_{\\rm H}}= \\frac{G M_\\star}{R_\\star} {\\rm ~~i.e.~~} {\\rm k}T_{\\rm g}\\approx 45 {\\rm ~MeV~}(M_\\star/{\\rm M}_\\odot) (R_\\star/ 10{\\rm ~km})^{-1} \\label{tg} \\end{equation} assuming ionised hydrogen ($\\mu=0.5$). This is twice the virial temperature of bound particles in circular orbit at $R_{\\star}$. Accretion onto a compact object can power the emission of gamma rays.  \\subsection{Disc formation} Matter infalling onto the compact object will generally have some non-zero angular momentum. A particle with a ballistic trajectory will therefore have $R^2\\Omega=(R^2\\Omega)_{{\\rm t}=0}$. At the radius of closest approach to the compact object $1/2 (R \\Omega)_{\\rm c}^2 = GM/R_{\\rm c}$ so that $2GMR_{\\rm c}=(R^2\\Omega)_0^2$. Only a particle with very low angular momentum can directly hit the compact object ($R_{\\rm c}<R_\\star$). This condition is not met in compact binaries where matter comes from an orbiting companion and the stream will go round the compact star and intersect itself. Subsequent shocks lead to the dispersion of energy and the stream settles onto a circular orbit with the initial angular momentum (see \\citealt{lubow} for details).  A steady supply of matter with some specific initial angular momentum piles up in a ring at this circularisation radius. There is no accretion unless some matter can transfer its angular momentum to reach smaller orbits. Under such a process, an accretion disc forms extending down to the compact object.  \\subsection{Angular momentum transport}  One process by which particles may exchange angular momentum is viscosity (see e.g. \\citealt{terquem} for an introduction). In gas kinetic theory, molecular viscosity arises from the exchange of momentum across the surface of a fluid element. The resulting force is proportional to $\\nu_{\\rm mol}\\sim \\lambda u$ where $\\lambda$ is the mean free path between collisions and $u$ is the mean thermal speed of the particles. Assuming particles in a plasma interacting only via Coulomb forces leads to: \\begin{eqnarray*} & \\lambda \\sim & 1~{\\rm cm}~(T/10^5~{\\rm K})~(\\rho/10^{-8}~{\\rm g~cm}^{-3}) \\nonumber \\\\ & u \\sim & 10^6{\\rm ~cm~s}^{-1}~(T/10^5{\\rm ~K})^{1/2} \\end{eqnarray*} using typical accretion disc temperatures and densities \\citep{fkr}. The viscosity coefficient is therefore $\\nu_{\\rm mol}\\sim 10^6{\\rm ~cm^2~s^{-1}}$.  Anticipating a little bit on result from the thin disc model, we can get an order-of-magnitude observational estimate for viscosity in accretion discs. If the eruptions of dwarf novae, a sub-class of CVs (see \\S\\ref{cv}), are due to the accretion of matter coming from the outer regions of a disc and transported by viscosity then $\\nu_{\\rm disc}\\sim R_{\\rm d} v_{\\rm r}$ where $R_{\\rm d}$ is the outer disc radius and $v_{\\rm r}$ is the radial velocity of the infalling matter. $v_{\\rm r}\\sim R_{\\rm d}/\\tau_{\\rm e}$ where $\\tau_{\\rm e}$ is the timescale of the eruption. The disc radius can be estimated using e.g. the circularisation radius (\\S1.3) and typically will be $\\sim 10^{10}$~cm. For an eruption lasting a day, $\\nu_{\\rm disc}\\sim 10^{15}{\\rm ~cm^2~s^{-1}}$ i.e. $\\nu_{\\rm disc}\\gg \\nu_{\\rm mol}$. Molecular viscosity is much too weak to account for accretion in discs.  The process by which angular momentum is transported has been the subject of intense research with turbulent transport the prime suspect. Turbulent viscosity is often modelled using the Navier-Stokes formalism but with a coefficient $\\nu_{\\rm turb}\\sim \\lambda_{\\rm turb} u_{\\rm turb}$ where the scale and speed are those of the turbulent eddies. These are dynamic properties of the fluid, not intrinsic as with molecular viscosity. In a disc, the size and speed of the eddies are likely to be limited by the scale-height $H$ and sound speed $c_{\\rm s}$: \\begin{equation} \\nu=\\alpha c_{\\rm s} H \\label{ssp} \\end{equation} where $\\alpha$ is a parameter $<1$. This is the famous $\\alpha$ parameterisation of disc viscosity first described in a landmark paper by \\cite{ss}. In the following I will use this parameterisation.  One would rather want to derive $\\nu$ from first principles. At present, only turbulence arising from the magneto-rotational instability (MRI) has succeeded in predicting any significant viscosity (see \\citealt{bh} for a review). Fortunately, there is some sense in modelling the angular momentum transport and energy dissipation of MHD turbulence using the $\\alpha$ parameterisation \\citep{bp}.  However, the mechanism may not work in weakly ionised flows and this is a problem for models with a cold accretion disc \\citep{stepinski,menou}. Putative global hydrodynamical instabilities might take over in these conditions. There are other possibilities for angular momentum transport including instabilities in a self-gravitating disc (e.g. \\citealt{bp} and references therein), spiral waves excited by tidal torques \\citep{spruit}, hydromagnetic winds launched from the disc \\citep{blapay}, radiative viscosity \\citep[e.g.][]{loeb} etc. However, these only apply in specific conditions.  \\subsection{Vertically integrated disc equations} The disc equations derive from the equations of fluid dynamics combined with a model for viscosity (which contains the turbulent magnetic field contributions in the MRI case), an equation of state for the gas and a description of the radiative processes. A standard set of assumptions to start with is: \\begin{itemize} \\item Axisymetry so that $\\partial/\\partial\\phi=0$ in cylindrical coordinates. \\item The only non-zero component of the stress tensor is the azimuthal shear $\\tau_{r\\phi}$. \\item The gas is perfect so the internal energy per unit mass is $e=c_vT$ and the gas pressure $P=\\rho e (\\gamma-1)$ with $\\gamma=c_p/c_v$. Radiation pressure is neglected here for simplicity. \\end{itemize} In addition I suppose the disc self-gravity is negligible (this assumption must be dropped for protostellar discs and AGNs) and that relativistic corrections are negligible (which is fine when more than a few gravitational radii away from the compact object). Assuming hydrostatic balance so that $v_{\\rm z}=0$, the vertical momentum conservation is ($P=\\rho c_s^2$): \\begin{equation} \\pd{P}{z} = -\\rho g_z {\\rm ~~i.e.~~} \\pd{\\ln P}{\\ln z}=-\\frac{\\Omega^2_{\\rm K}}{c_s^2}~ z^2 \\left[1+\\frac{z^2}{R^2}\\right]^{-\\frac{3}{2}} \\label{hydro} \\end{equation} If the height $H$ from the midplane of the disc is $\\lta R$ high order terms in $z/R$ can be neglected. This equation can be integrated analytically for a perfect gas yielding $P(z)$ and $\\rho(z)$. Averaging $P$ over $z$ gives a relationship between $H$ and the mid-plane sound speed which, to a factor of order unity, is : \\begin{equation} H= c_s/\\Omega_{\\rm K} \\end{equation} The scale-height $c_s/\\Omega_{\\rm K}$ appears in Eq.~\\ref{hydro}. The detailed vertical balance can be bypassed by assuming the disc height is given by this relationship. The Shakura-Sunyaev prescription for the viscosity becomes \\begin{equation} \\nu=\\alpha c_s^2 / \\Omega_{\\rm K} \\end{equation} Note that $\\nu$ is effectively integrated over $z$. With a Navier-Stokes formulation of the viscosity, the integrated stress is: $$\\tau_{r\\phi}=\\nu \\Sigma R \\pd{\\Omega}{R}$$ where $\\Sigma=2 \\rho_o H$ is the column density and $\\rho_o$ is the mean density. The radial evolution equations are then integrated and decoupled from $z$, resulting in a set of time-dependent 1D equations: \\begin{eqnarray} \\pd{\\Sigma}{t}+\\frac{1}{R}\\pd{}{R}(\\Sigma R v_r) & = & 0 \\label{mass}\\\\ \\pd{v_r}{t}+v_r \\pd{v_r}{R} & = & R\\Omega^2-R\\Omega_{\\rm K}^2-\\frac{1}{\\rho_o}\\pd{P}{R} \\label{radial}\\\\ \\pd{R^2 \\Omega}{t}+v_r\\pd{R^2 \\Omega}{R} & = & \\frac{1}{\\Sigma R}\\pd{}{R}\\left(R^2\\tau_{r\\phi}\\right)  = \\frac{1}{\\Sigma R}\\pd{}{R}\\left(\\nu \\Sigma R^3\\pd{\\Omega}{R}\\right) \\label{angular}\\\\ T\\pd{s}{t}+v_r T\\pd{s}{R} & = & \\frac{1}{\\Sigma}(Q^+_{\\rm vis}-Q^-_{\\rm rad}) \\label{energy} \\end{eqnarray} In order, we have the mass, radial momentum, angular momentum, and energy conservation equations. $R\\Omega$ and $v_r$ are the angular and radial velocities while $P=\\rho_o c_s^2$. $c_s$ and therefore $T$ are the mid-plane values of the sound speed and temperature in this framework (Eq.~1.8). Again, only the dominant term in the radial component of gravity is kept $g_r\\approx R\\Omega_{\\rm K}$. In the thermal equation, $s$ is the entropy and  $Q^+_{\\rm vis}$ is the energy generated by the work of viscous forces in the annulus: \\begin{equation} Q^+_{\\rm vis}=\\nu \\Sigma \\left(R\\pd{\\Omega}{R}\\right)^2 \\label{qplus} \\end{equation} $Q^-_{\\rm rad}$ represents radiative losses.  Using $ds=de+Pd(1/\\rho_o)$ the thermal equation can be rewritten as: \\begin{equation} \\pd{e}{t}+v_r \\pd{e}{R} = -\\frac{P}{\\rho_o^2}\\left[ \\pd{\\rho_o}{t}+v_r \\pd{\\rho_o}{R}\\right] +\\frac{1}{\\Sigma} (Q^+_{\\rm vis}-Q^{-}_{\\rm rad}) \\end{equation} and $e$ can be written using $T_{\\rm g}$ (see Eq.~\\ref{tg}): \\begin{equation} e=\\frac{c_s^2}{(\\gamma-1)}=\\frac{T}{T_{\\rm g}} R^2 \\Omega_{\\rm K}^2 \\label{internal} \\end{equation} Other versions of these equations accommodate additional energy sources and sinks, conduction or convection terms, radiation pressure, different proton and electron temperatures, magnetic fields, relativistic corrections etc.   \\subsection{Steady state disc} A disc in steady state has $\\pd{}{t}=0$. Combining Eqs.~\\ref{mass} and \\ref{angular} and integrating over $R$ gives: \\begin{equation} R^2\\Omega=\\frac{\\nu R^2}{v_r}\\pd{\\Omega}{R}+{\\rm ~constant} \\label{int1} \\end{equation} Assuming there are no viscous torques at the surface of the compact object, the integration constant is $R_\\star^2 \\Omega_\\star$. This term is very small compared to $R^2\\Omega$ far from the boundary layer. Writing $b=1-(R^2_\\star \\Omega_\\star/R^2 \\Omega)$, Eq.~\\ref{int1} then gives \\begin{equation} b v_r= \\frac{\\nu}{R}\\pd{\\ln \\Omega}{\\ln R}\\label{vr} \\end{equation} The radial momentum equation can be rewritten using Eqs.~\\ref{vr} and \\ref{internal} as: \\begin{equation} \\frac{\\Omega^2}{\\Omega^2_{\\rm K}}-1=(\\gamma-1)\\left( \\frac{T}{T_{\\rm g}} \\right) \\pd{\\ln P}{\\ln R}+(\\gamma-1)^2\\left( \\frac{\\alpha}{b} \\frac{T}{T_{\\rm g}} \\right)^2 \\left( \\pd{\\ln\\Omega}{\\ln R} \\right)^2  \\pd{\\ln v_{r}}{\\ln R} \\label{radialadim} \\end{equation} Defining the cooling efficiency $f$ as $1-Q^-_{\\rm rad}/Q^+_{\\rm vis}$, the energy equation becomes: \\begin{equation} f b \\left(\\frac{\\Omega}{\\Omega_{\\rm K}} \\right)^2 \\pd{\\ln \\Omega}{\\ln R}=\\left( \\frac{T}{T_{\\rm g}} \\right) \\left[ \\pd{\\ln e}{\\ln R}-(\\gamma-1)\\pd{\\ln \\rho_o}{\\ln R} \\right] \\label{eneradim} \\end{equation} In a solution to the equations, $f$ should be calculated self-consistently from the radiative processes.  I find this set of dimensionless equations useful when discussing the assumptions that go into a thin or thick disc. ",
        "conclusions": ""
    },
    "0206/astro-ph0206510_arXiv.txt": {
        "abstract": "The  Fornax Cluster Spectroscopic Survey (FCSS) project utilises the Two-degree field (2dF) multi-object spectrograph on the Anglo-Australian Telescope. Its aim is to obtain spectra for a complete sample of all 14000 objects with $16.5\\leq b_j \\leq19.7$ {\\em irrespective of their morphology} in a 12~deg$^{2}$ area centred on the Fornax Cluster. A  sample of 24 Fornax Cluster members has been identified from the first 2dF field (3.1~deg$^{2}$ in area) to be completed. This is the first complete sample of cluster objects of known distance with well defined selection limits. 19 of the galaxies (with --15.8$<M_{B}<$--12.7) appear to be conventional dwarf elliptical (dE) or dwarf S0 (dS0) galaxies. The other 5 objects (with  --13.6$<M_{B}<$--11.3) are those galaxies which we described  in Drinkwater  et al. (2000b) and labelled `Ultra-Compact Dwarfs' (UCDs). A major result is that the conventional dwarfs  all have   scale-sizes $\\alpha$$\\ga$3~arcsec ($\\simeq$300~pc). This apparent minimum scale size implies an equivalent minimum luminosity  for a  dwarf of a given surface brightness. This produces a  limit on their distribution in the  magnitude-surface brightness plane, such that we do not observe dEs with high surface-brightnesses but  faint absolute magnitudes.  Above this observed minimum scale-size  of 3~arcsec, the dEs and dS0s fill the whole area of the magnitude--surface brightness plane sampled by our selection limits.  The observed correlation between magnitude and surface brightness noted by several recent studies of brighter galaxies is not seen with our  fainter cluster sample. A comparison of our results with the Fornax Cluster Catalog (FCC) of Ferguson illustrates that attempts to determine cluster membership solely on the basis of observed morphology can produce significant errors. The FCC  identified 17 of the 24 FCSS sample (i.e 71 per cent) as being `cluster' members, in particular  missing all 5 of the UCDs. The FCC also suffers from significant contamination:  within the FCSS's field and selection limits, 23 per cent of those objects described as  cluster members by the FCC are shown by the FCSS to be background objects. ",
        "introduction": "It is generally perceived that dwarf galaxies are primarily of low surface brightness, and that  lower luminosity galaxies have fainter surface brightnesses. In particular, low surface brightness galaxies (LSBGs) seen towards a cluster are conventionally assumed to be members, while apparently faint, but high surface brightness galaxies (HSBGs) are presumed to be luminous objects in the background (e.g. Sandage, Binggeli \\& Tammann 1985). Indeed, it has frequently been asserted that dwarf galaxies (in particular) exhibit a strong surface brightness -- luminosity relation (e.g. Sandage, Binggeli \\& Tammann 1985; Bingelli, Sandage \\& Tammann 1985; Kormendy 1985; Ferguson \\& Sandage 1988). Binggeli (1994) has particularly persuasively argued this point with reference to the Local Group. Nevertheless, other authors have noted that a good deal of the evidence for the correlation relies on the `eyeball' selection of dwarf galaxies `similar' to ones already known, that any correlation breaks down at the low luminosity end, and/or that there may be strong selection effects limiting the detected galaxies' photometric parameters (e.g. Phillipps, Davies \\& Disney 1988; Irwin et al. 1990).  The failure of the relation, i.e. the existence of large background LSBGs (such as the serendipitously discovered Malin 1; Bothun et al.\\ 1987), or of a population of high surface brightness (compact) dwarfs in a cluster (Drinkwater \\& Gregg 1998), or indeed of a population of galaxies so compact as to masquerade as stars and, hence, be missed altogether from galaxy samples (e.g. Arp 1965; Drinkwater et al. 2000b; Phillipps et al. 2001) could have a dramatic effect on our perception of the galaxy population as a whole (see Impey \\& Bothun 1997).  The primary factor preventing a resolution of the controversy has been the non-existence of suitable samples of confirmed dwarfs of known distance (and hence luminosity) with well defined selection limits. Generally, dwarf galaxy samples have been either subjectively (for instance, morphologically) selected sets of cluster galaxies (e.g. Ferguson 1989) or objectively chosen sets of galaxies with unconfirmed cluster membership/unknown distances (e.g. Irwin et al. 1990).   However, the development of a new generation of multi-object spectrographs, exemplified by the `2 degree field', or 2dF, multi-fibre spectrograph on the Anglo-Australian Telescope (AAT) (see e.g. Taylor, Canon \\& Parker 1998), has  made it possible to conduct a truly complete spectroscopic survey of a given area on the sky, down to well determined, faint limits, irrespective of image morphology or any other pre-selection of target type, and thus obtain redshifts for a large and unbiased sample of targets (Phillipps 1997). Our  Fornax Cluster Spectroscopic Survey (FCSS) aims to  exploit the huge multiplexing advantage of 2dF in this way by surveying, in total, a region of 12~deg$^{2}$ centred on the Fornax Cluster of galaxies. The prime reason for choosing this area of  sky area was, of course, the presence of the Fornax Cluster itself. Fornax, with Virgo, is the nearest reasonably rich cluster (approximately Abell Richness Class 0 -- it is supplementary cluster S0373 in Abell, Corwin \\& Olowin 1989).    The targets of the FCSS encompass both cluster galaxies, of a wide range of types and magnitudes, and background galaxies (over a similarly wide range of morphologies), as well as Galactic stars and QSOs (Drinkwater et al. 2000a). Our main motivation for an all-object survey was to determine cluster membership for as complete a sample of objects as possible across a wide range of surface brightnesses. We wished to test, in particular, whether the usual assignment of LSBGs to the cluster and high (or even `normal') surface brightness faint galaxies to the background is justified. The situation regarding HSBGs is even less clear than for LSBGs -- the question of the very existence of normal or high surface brightness dwarf galaxies (excepting possibly pathological cases like M32) has been highly controversial -- see, for example, the contrasting views expressed in Ferguson \\& Sandage (1988) and Irwin et al.\\ (1990). An up to date discussion is given in Trentham \\& Hodgkin (2002).   Drinkwater et al. (2000a) (hereafter Paper~I) outlined the aims and methods of the FCSS and presented the results from the first  2\\degr\\, diameter field surveyed. Drinkwater et al. (1999) (Paper~II)  described the nature of several compact galaxies found beyond the cluster. Meyer et al. (2001) ( Paper~IV) discussed the quasars found in the FCSS data. Drinkwater et al. (2000b) (Paper~III) and Phillipps et al. (2001) (Paper~V) described the discovery of 5 extremely compact, high surface brightness dwarf galaxies within the cluster. These objects (dubbed `Ultra-Compact Dwarfs' -UCDs) have intrinsic sizes $\\la$100~pc. They are more compact and significantly less luminous than other known compact dwarf galaxies, yet are much brighter than globular clusters.    The data from the first field (see Paper~I) revealed that the Fornax Cluster is clearly defined kinematically. A total of 24 galaxies from the data lie within the cluster (this includes the 5 UCDs described in Papers~III and V). This is, to our knowledge, the first complete sample of dwarf galaxies of known distance (and hence luminosity) with well defined selection limits. In this paper we study the nature of the objects in this complete sample. Section 2 briefly describes the methods and present status of the FCSS. Section~3 describes the contents of the complete Fornax Cluster sample and presents  $B$ and $R$ photometry, radial intensity profiles and radial colour distributions for all 24 confirmed cluster members. Using this data we are able to explore (in Section 4) the distribution in luminosity and surface brightness of the confirmed cluster members. Section 5 presents our conclusions. ",
        "conclusions": "We have obtained spectra for a morphologically unbiased set of cluster dwarf galaxies as part of the  Fornax Cluster Spectroscopic Survey. The present paper provides the photometric study of these confirmed cluster dwarfs. We find that, despite their low luminosities, not all the dwarfs also have low surface brightness. There is therefore a wide spread of parameters in the space of surface brightness versus luminosity or scale size. The only correlation evident in these plots is the one in the luminosity--surface brightness plane caused by the absence of `normal' dwarfs with scale sizes below about 300~pc. Even then, a separate class of Ultra-Compact Dwarfs fills some of this otherwise empty area of parameter space. Morphological selection of cluster dwarfs solely on the grounds of their low surface brightness may therefore overlook a significant number of dwarf galaxies which are virtually impossible to discriminate from background, larger galaxies on appearance alone.   The data presented within this paper result from the completion of only one of the intended 4 FCSS fields.  We now have full data for the second field and are in the process of reducing these. The completion of all 4 fields will not only improve the statistics of our results  but will also enable us to study  the distribution of dwarf galaxies of different morphologies (i.e. dE, dI, UCD) as a function of position within the cluster. Such studies will be assisted by the deep  CCD g,r,i,z  imaging of the cluster region which we have now obtained using the 4.0-m Blanco Telescope  at the Cerro Tololo Inter-American Observatory. These data are currently  being reduced.   Important  scientific advances  could also potentially be made by improving the sensitivity of our spectroscopy to fainter luminosities and/or surface brightnesses. This can be clearly seen by studying Figs.~4 and 5. Using the  CCD data we should be able to define a sample to a considerably fainter $B$ magnitude limit.  The standard 2dF set-up would be adequate to obtain spectra of the higher surface brightness objects in such a sample.  This would include many potential UCDs. However, even bigger advances will require us to obtain spectra for objects at fainter surface brightness limits than the present limit of 23.7~B~mag~arcsec$^{-2}$. This possibility now exists with the commissioning of multi-object spectrographs on the new generation of 8-m telescopes.   The number of cluster objects found in Field 1 is dwarfed by the number of confirmed background objects (1175). Such a sample is potentially of great value to studies of the BBD of field galaxies. Although the present magnitude limit is only slightly lower than that of the 2dFGRS used by Cross  et al. (2001) for their determination of the BBD,  the FCSS background sample has the advantage that  the FCSS sources were not pre-selected as galaxies. Hence, our sample will  contain a fairer representation of compact objects in the background than Cross et al.'s (see Paper~II). For the lower surface brightness objects in our sample we  have also used longer exposure times than the 2dFGRS, so FCSS extends to fainter surface brightness than 2dFGRS. However, if the FCSS sample could be pushed to even fainter luminosity and surface brightness limits then, not only will this enable us to study the cluster BBD better, but the resulting sample will be ideal for studying the field BBD to fainter surface brightness limits than  yet obtained.  For example, Cross et al.'s correlation between luminosity and surface brightness has only been tested between  --24$<M_{B}<$--16.0 and 18.0$<\\mu_{e}<$24.5. At present the FSCC is  $\\simeq$80 per cent complete to a surface brightness about 0.5~mag fainter than this. Further improvements in the sensitivity of the spectroscopy will enable us to sample whole new regions of ($M_{B}$,$\\mu_{o}$)  space."
    },
    "0206/astro-ph0206383_arXiv.txt": {
        "abstract": "We explore several models which might be proposed to explain recent possible detections of high-energy (TeV) gamma rays in association with low-energy gamma-ray bursts (GRBs). Likely values (and/or upper limits) for the source energies in low- and high-energy gamma rays and hadrons are deduced for the burst sources associated with possible TeV gamma-ray detections by the Project {\\it GRAND} array.  Possible spectra for energetic gammas are deduced for three models: 1) inverse-Compton scattering of ambient photons from relativistic electrons; 2) proton-synchrotron emission; and 3) inelastic scattering of relativistic protons from ambient photons creating high-energy neutral pions, which decay into high-energy photons.  These models rely on some basic assumptions about the GRB properties, e.g. that:  the low- and high-energy gamma rays are produced at the same location; the time variability of the high-energy component can be estimated from the FWHM of the highest peak in the low-energy gamma ray light curve; and the variability-luminosity relation of Fenimore \\& Ramirez-Ruiz (2000) gives a reliable estimate of the redshifts of these bursts. We also explore the impact of each of these assumptions upon our models. We conclude that the energetic requirements are difficult to satisfy for any of these models unless, perhaps, either the photon beaming angle is much narrower for the high-energy component than for the low-energy GRB or the bursts occur at very low redshifts ($\\lesssim 0.01$). Nevertheless, we find that the energetic requirements are most easily satisfied if TeV gamma rays are produced predominantly by inverse-Compton scattering with a magnetic field strength well below equipartition or by proton-synchrotron emission with a magnetic field strength near equipartition. ",
        "introduction": "\\label{sec:level1}  Evidence has been accumulating for the arrival of $\\sim$ GeV-TeV gamma rays in coincidence with low-energy ($\\sim $ MeV) gamma-ray bursts (GRBs). For example, {\\it EGRET} detected seven GRBs which emitted high energy photons in the $\\sim 100$ MeV to 18 GeV range \\citep{Schneid,  Hurley, Catelli}. There have also been results suggestive of gamma rays beyond the GeV range \\citep{Amenomori, Padilla}, although these were not claimed as firm detections. Evidence  has also been reported for TeV emission in one burst out of 54 {\\it BATSE} GRBs in the field of view of the {\\it Milagrito} detector \\citep{milagro}.  In another paper \\citet{Poirier} reported suggestive evidence for sub-TeV gamma rays arriving in coincidence with GRBs which occurred near zenith above the Gamma Ray Astrophysics at Notre Dame ({\\it GRAND}) air shower array.   In that experiment, most of the eight bursts analyzed were associated with at least some marginal excess ($\\sim  1\\sigma$) of muons including the event detected by {\\it Milagrito}.  One burst evidenced a possible detection at the $2.7\\sigma$ level. As shown in \\citet{Poirier},  if this detection is real, then the output in energetic gammas is likely to dominate the energetics of the burst.  Although these data are not overwhelmingly convincing, they are at least suggestive that detectable energetic TeV gamma rays might be associated with low-energy gamma-ray bursts \\citep{Vernetto}. Moreover, these new detections, if real, can not be explained by a simple extrapolation of the {\\it BATSE} spectrum \\citep{Poirier}, particularly if intergalactic absorption is taken into account \\citep{Salamon, tot00}. A new $\\sim$ TeV component in the GRB spectrum seems to be required.   The present work is therefore  an attempt to interpret these new possible detections in the context of  three models, which might be proposed for the production of TeV gammas in a GRB.  These are: 1) inverse-Compton scattering of ambient photons from relativistic electrons in the burst environment; 2) proton-synchrotron emission \\citep{Vietri97, tot98a, tot98b, tot00}; and 3) inelastic scattering of relativistic protons from ambient photons creating high-energy neutral pions, which decay into high-energy photons \\citep{Waxman95, Waxman97}. We here briefly outline the underlying physics and characteristic energetic gamma-ray spectra associated with each of these possible models.  We then derive limits on the parameters of these models based upon the detection limits from the Project {\\it GRAND} array.  Based upon this, we conclude that it is difficult for any of these models to satisfy the energetic requirements unless the photon beaming angle is very narrow for the high-energy component. Of the models considered, the most likely are inverse-Compton scattering or proton-synchrotron emission.  We note, however, that these conclusions rely upon a few assumptions.  For example, we have assumed that the low- and high-energy gamma rays are produced at the same location; that the time variability of the high-energy component can be estimated from the FWHM of the highest peak in the low-energy gamma ray light curve; and that the variability-luminosity relation of Fenimore \\& Ramirez-Ruiz (2000) gives a reliable estimate of the redshifts of these bursts.  We explore the impact of each of these assumptions and find that, unless the bursts occur at very low redshifts ($\\lesssim 0.01$), the energetic requirements remain difficult to satisfy. ",
        "conclusions": "We have analyzed the eight GRBs discussed in \\citet{Poirier}, which occurred above the Project {\\it GRAND} array. We have studied the implied energies in energetic ($\\sim$TeV) gamma rays (and associated electrons and protons) of such GRBs in the context of three possible mechanisms: inverse-Compton scattering, proton-synchrotron radiation, and photo-pion production.  Our analysis suggest that all of these models face significant energetic requirements.  Gamma-ray production by either inverse-Compton scattering or proton-synchrotron radiation is probably the most efficient process.  Although it can not be claimed that TeV gammas have unambiguously been detected in association with low-energy GRBs, we have argued that there is enough mounting evidence to warrant further study.  Furthermore, we have shown that if TeV gammas continue to be observed, then they present some interesting dilemmas for GRB physics. In view of their potential as a probe of the GRB source environment, we argue that further efforts to measure energetic gamma rays in association with low-energy gamma-ray bursts are warranted."
    },
    "0206/hep-ph0206163_arXiv.txt": {
        "abstract": "\\noindent We review some astrophysical and cosmological properties and implications of neutrino masses and mixing angles. These include: constraints based on the relic density of neutrinos, limits on their masses and lifetimes, BBN limits on mass parameters, neutrinos and supernovae, and neutrinos and high energy cosmic rays. ",
        "introduction": "The role of neutrinos in cosmology and astrophysics can not be understated \\cite{dolgovreview}.  They play a critical role in the physics of the early Universe, at temperatures scales of order 1 MeV, and strongly determine the abundances of the light elements produced in big bang nucleosynthesis. They almost certainly play a key role in supernova explosions, and if they have mass, could easily contribute to the overall mass density of the Universe. At the present time, the only indicators of neutrino masses are from astrophysical sources, the inferred neutrino oscillations of neutrinos produced in the Sun, and those produced in cosmic-ray collisions in the atmosphere. Indeed, their elusive character has proven that a great deal of information on neutrino properties can be gained by studying their behavior in astrophsyical and cosmological environments. Here, we will try to elucidate some of these constraints on neutrino masses.     In our discussion below, we will assume that the early Universe is well described by a standard Friedmann-Lemaitre-Robertson-Walker metric \\beq ds^2 = dt^2 - R^2(t) \\left [ \\frac{d r ^2}{1 - k r ^2} + r^2\\, (d \\theta^2 + \\sin^2\\theta\\, d \\phi^2 ) \\right ] \\eeq We further assume that thermal equilibrium was established at some early epoch and that we can describe the radiation by a black body equation of state, $p = \\rho/3$ at a temperature $T$. Solutions to Einstein's equations allow one to determine the expansion rate of the Universe defined to be the Hubble parameter in terms of the energy density in radiation, the curvature and the cosmological constant. In the early Universe the latter two quantities can be neglected and we write \\beq H^2 \\equiv \\left ( \\frac{\\dot R}{R} \\right )^2 = \\frac{8 \\pi\\; G_{N}\\; \\rho}{3} \\label{H} \\eeq where the energy density is \\beq \\rho {} = \\left( \\sum_B g_{B} + {7 \\over 8} \\sum_F  g_{F} \\right) {\\pi^{ 2} \\over 30}  T^{4} \\equiv    {\\pi^{ 2} \\over 30}\\, N(T) \\, T^{4} \\label{rho} \\eeq The present neutrino contribution to the total energy density, relative to  the critical density (for a spatially flat Universe) is \\beq \\Omega_\\nu = \\frac{\\rho_\\nu}{\\rho_c} \\label{omegaratio} \\eeq where $\\rho_c  = 1.06 \\times 10^{-5} {h}^2 {\\rm GeV/cm}^3$ and $h = H/100$km/Mpc/s is the scaled Hubble parameter. For a recent review of standard big bang cosmology, see~\\cite{OP}. ",
        "conclusions": ""
    },
    "0206/astro-ph0206330_arXiv.txt": {
        "abstract": "\\noindent In this paper, timing data for all of the anomalous X-ray pulsars and soft gamma repeaters are compiled. Timing properties of these objects are investigated. The effect of bursts of soft gamma repeaters on their period history is investigated. The P-\\.{P} diagram for pulsars, X-ray binaries, anomalous X-ray pulsars, soft gamma repeaters and dim radio quiet netron stars is constructed. The possible evolutionary tracks for anomalous X-ray pulsars, soft gamma repeaters and dim radio quiet netron stars are examined. \\\\ \\\\ KEY WORDS: AXP, SGR, DRQNS, MAGNETAR ",
        "introduction": "The new kind of neutron stars (NS); so-called anomalous X-ray pulsars (AXPs) and soft gamma repeaters (SGRs) are characterized by their rapid spin down $[1]$. To explain this rapid spin-down, and the other characteristics of these objects such as the low rotational energy loss, two fundamentally different models are proposed: accretion from a fall-back disk model and the magnetar model. Accretion from a fall-back disk model describes these objects as systems with disks which have just finished the so-called propeller phase $[2]$. However, it is shown that accretion from a fall-back disk model has too many difficulties in the explanation of AXPs $[3]$. An expansion of this model might include the effect of propeller and accretion at the same time. Such a model will be considerably successful in the explanation of properties of AXPs, but not SGRs $[3]$.   To explain the strong gamma-ray bursts and the persistent X-ray radiation, it is proposed that the magnetic field of SGRs and AXPs decay in $\\sim$10$^4$-10$^5$ yr. Moreover, the P-\\.{P} values and the strength of the giant gamma ray bursts suggest very high magnetic fields for these objects. According to the magnetar model, AXPs and SGRs have magnetic fields of $10^{14}-10^{15}$ G $[4]$. Magnetar models are more attractive because they utilize new physical phenomena trying to explain the AXP and SGR properties together simultaneously. Since accretion theory is very well developed and well known, it is much easier to criticize the accretion models proposed for AXPs and SGRs. This is the first of a series of papers showing the problematic and good aspects of the magnetar model depending on the observational data.   Three AXPs are connected to supernova remnants (SNRs) $[5,6]$ and one of the SGRs, namely SGR 0526-66, is possibly genetically connected to SNR N49 in the Large Magellanic Cloud [6]. About 21 radio pulsars (PSRs) are connected to galactic SNRs $[7,8,9,10]$ with high probability. Two of these AXPs (AXP 1E1841-045 and AXP 1E2259+586 which are located in SNRs G27.4+0.0 and G109.1-1.0 respectively) are not farther than 7 kpc from the Sun. Observational data show that in that volume the number of SNRs is not less than 100. Contrary to PSRs, AXPs can be detected at such distances relatively easily. From the above arguments, it follows that the birth rate of AXPs is at least 50 times smaller than the birth rate of SNRs.  Some recently found PSRs have magnetic fields of B=$10^{13}-10^{14}$ G  $[11]$. We do not yet know the properties of these PSRs in detail, but young NSs with B$<10^{14}$ G (having ages less than 10$^5$ years) in general, have very low X-ray luminosities and do not have $\\gamma$ ray burst activity. However, observed NS-magnetars always have L$_x > 10^{34}$ erg/s and they sometimes have strong $\\gamma$ ray burst properties. Whether this statement about NS-magnetars is true in general is not known since number of data are few.  AXPs and SGRs are distant objects. Therefore it is very difficult to obtain trustworthy and sensitive enough observational data from these objects. This is also true for the SNRs which might be genetically connected with them. Moreover, the number of AXPs and SGRs are so small to perform statistical analysis. This situation cannot considerably change in the future. The direct estimate of the magnetic field strength of these objects from cyclotron line detection is more important. It is such a pity that we do not yet detected such a cyclotron line. Therefore for the investigation of origin and physical nature of these very important objects, it is necessary to study also neutron stars which demonstrate other types of properties. Young radio pulsars (PSRs) and single radio quiet neutron stars (DRQNS) which are observed as dim point X-ray sources will provide valuable information. In this paper we concentrate our attention on the timing properties of AXPs and SGRs. Also, we want to understand the possible evolution of AXPs and SGRs on the P-\\.{P} diagram, qualitatively. ",
        "conclusions": "1. The absolute deviations of period values from the constant \\.{P} line in period histories are larger for SGR1900+14, SGR1806-20 and AXP1048.1-5937 which have the largest values of \\.{E} in comparison to all other AXPs. In the frame of the magnetar model, these objects must also have large values of temperature and persistent luminosities since they have higher magnetic field decay. However, these sources do not have the largest temperature and luminosity among AXPs on the average. This is hard to understand in the frame of the magnetar model. Moreover, we expect to see the same behavior of unobserved bursts and noise from AXPs 1E1841-045 and RXJ170849-4009 similar to AXP 1048.1-5937, as the number of period data increase. If burst and noise characteristics are not observed from these AXPs, then this would be a handicap for the magnetar model.  2. We estimated the lower limit of \\.{P} value of AXP J1845.0-0300 to be about 1.3 10$^{-12}$ s/s using the SNR age as a constraint. The existence of magnetic decay decreases the upper limit of \\.{P}. We conclude that \\.{P} should be in the range between (1.3-2) 10$^{-12}$ s/s.  3. The large increase in the P and \\.{P} values of SGR 1900+14 and SGR 1806-20 after the burst activity periods might be caused by plasma ejection and interaction with the surrounding matter. The time-scale of increase of P which is about 1-3 years for these two objects, shows that the process increasing the period finishes its effect when the ejected plasma reaches less than 1 pc from the NS. The high P values of these SGRs are expected to fall to lower values in the future.  4. SGRs and AXPs are not sources having strong PWN in spite of the existence of bursts for SGRs and possible burst activity for AXPs. It is shown that PWN appears only if pulsar have small ages ($<$10$^4$ years) and simultaneously very large value of \\.{E} (which is proportional to voltage). Nearly all PSRs with $\\tau < 1.6\\times10^4$ years up to distance 5 kpc were found to be associated with SNRs. All 21 PSRs which have connection with SNRs, have values of \\.{E}$>$3$\\times10^{34}$ erg/s. But PWN have only been observed from 8 PSRs which have $\\tau < 3\\times 10^4$ years and \\.{E}$> 4\\times10^{35}$ erg/s. On the other hand, near the SGR1900+14, SGR1806-20, AXP1048.1-5937 and AXPJ170849-4009 with characteristic ages $<$10$^4$ years and distances d$<$10-15 kpc even no S-type SNRs are found. Evidently, it is not easy to find SNRs at such distances in the galactic center directions. But absence of PWN may be natural because for these objects \\.{E}$<$4$\\times 10^{34}$ erg/s.  5. Magnetar model considered SGRs and AXPs as one group of objects, so that AXP 1E2259+586 must have lost burst activity and underwent considerable magnetic decay. If we equate the characteristic age of this AXP to the age of SNR CTB 109, it is necessary for the initial magnetic field not to be smaller than the magnetic field of SGRs. From this fact and the location of SGRs and AXPs on the P-\\.{P} diagram (Figure 8) roughly on a hypothetical vertical line (P=const.) we may talk about strong magnetic field decay and large value of braking index for AXP/SGR group. It is also necessary to note that if source of burst energy is the magnetic field decay and because several DRQNSs are in the continuation of the evolutionary tracks of AXPs/SGRs, then the decay of magnetic field of AXPs and SGRs must continue further. If the DRQNSs with large periods, are the continuation of AXPs/SGRs, then we expect to observe noise from these pulsars higher than the noise of radio PSRs. Also, we expect to observe high noise from newly found PSRs which have magnetic fields closer to AXP 1E2259+586.  6. In Thompson \\& Duncan [4] model of magnetar, the order of magnitude of the magnetic field is estimated from P-\\.{P} values and also from the strength of the giant bursts. If these sources are involved in several bursts before, then they need to have several times higher magnetic field strength at birth than the current magnetic field strength or their magnetic field should be fed continuously by dynamo action. It is a possibility that the NSs with SGR and AXP properties might have lower mass (so larger in radius) and B$\\sim$10$^{13}$-10$^{14}$ G) to have burst activity and such noise. This makes the drift of monopoles much  easier, characteristic time for turbulence increases and consequently dynamo action works longer and magnetic field is fed longer. The NS with low mass perform pulsations for a long time until it comes to stability. This makes them cool later and increase the turbulence activity which feeds the dynamo.  7. Continuous increase of magnetic field for PSRs do not lead to AXP properties continuously. If in addition to the high magnetic fields we assume that the NS has lower mass, then, this causes \\.{P} to decrease sharply when NS reaches stability. Consequently the evolution from SGR to AXP and then to DRQNS do not require strong magnetic decay. In this case, we can say that there is no fundamental difference between PSRs and magnetars.                       $Acknowledgments$ We thank T\\\"{U}B\\.{I}TAK, the Scientific and Technical Research Council of Turkey, for support through TBAG-\\c{C}G4. We also would like to thank M. \\\"{O}zg\\\"{u}r Ta\\c{s}k\\i n for comments on the manuscript.   \\clearpage"
    },
    "0206/astro-ph0206276_arXiv.txt": {
        "abstract": "Wide-field photometric data from the Sloan Digital Sky Survey have recently revealed that the Galactic globular cluster Palomar 5 is in the process of being tidally disrupted (Odenkirchen et al.\\ 2001). Here we investigate the kinematics of this sparse remote star cluster using high resolution spectra from the Very Large Telescope (VLT). Twenty candidate cluster giants located within 6 arcmin of the cluster center have been observed with the UV-Visual Echelle Spectrograph (UVES) on VLT-UT2. The spectra provide radial velocities with a typical accuracy of 0.15~\\kms. We find that the sample contains 17 certain cluster members with very coherent kinematics, two unrelated field dwarfs, and one giant with a deviant velocity, which is most likely a cluster binary showing fast orbital motion. From the confirmed members we determine the heliocentric velocity of the cluster as $-58.7\\pm 0.2$~\\kms. The total line-of-sight velocity dispersion of the cluster stars is $1.1\\pm 0.2$~\\kms\\ (all members) or $0.9\\pm 0.2$~\\kms\\ (stars on the red giant branch only). This is the lowest velocity dispersion that has so far been measured for a stellar system classified as a globular cluster. The shape of the velocity distribution suggests that there is a significant contribution from orbital motions of binaries and that the dynamical part of the velocity dispersion is therefore still substantially smaller than the total dispersion. Comparing the observations to the results of Monte Carlo simulations of binaries we find that the frequency of binaries in Pal\\,5 is most likely between 24\\% and 63\\% and that the dynamical line-of-sight velocity dispersion of the cluster must be smaller than 0.7~\\kms (90\\% confidence upper limit). The most probable values of the dynamical dispersion lie in the range $0.12 \\le \\sigma/\\kms \\le 0.41$ (68\\% confidence). Pal\\,5 thus turns out to be a dynamically very cold system. Our results are compatible with an equilibrium system. We find that the luminosity of the cluster implies a total mass of only 4.5 to $6.0 \\times 10^3 M_\\odot$. We further show that a dynamical line-of-sight velocity dispersion of 0.32 to 0.37~\\kms\\ admits a King model that fits the observed surface density profile of Pal\\,5 (with $W_0 = 2.9$ and $r_t = 16\\farcm1$) and its mass. ",
        "introduction": "The globular cluster Palomar 5, an old halo cluster located at a distance of about 23~kpc from the Sun and 18.5~kpc from the Galactic center (see Harris 1996), stands out through a number of unusual and extreme properties. These make it particularly interesting from the viewpoint of dynamics, cluster evolution, and galactic structure. First of all, Pal\\,5 is extraordinarily sparse and faint. Its total luminosity of M$_V \\simeq -5.0$ (Sandage \\& Hartwick 1977, hereafter SH77) places it among the least luminous globular clusters that we currently know to exist in our Galaxy. Assuming a mass-to-light ratio typical of other globular clusters this luminosity corresponds to a stellar mass of about $1.3 \\times 10^4$~M$_\\odot$ (SH77) or $0.8 \\times 10^4$~M$_\\odot$ (Mandushev, Spassova \\& Staneva 1991), which lies a factor of 30 below the median mass of Galactic globulars. Further, Pal\\,5 has a very extended core and a low central concentration (see, e.g., Trager, King \\& Djorgovski 1995). Finally, the faint part of the luminosity function of Pal\\,5 is unusually flat (Smith et al.\\ 1986, Grillmair \\& Smith 2001), i.e.\\ the fraction of low-mass stars in Pal\\,5 is much smaller than in other galactic globulars. These peculiarities suggest that Pal\\,5 has undergone strong dynamical evolution and mass loss and that it may be close to complete disruption. The hypothesis of ongoing dissolution was recently confirmed by the discovery of two massive tails of unbound former cluster members which spread from the cluster in opposite directions over an arc of several degrees (Odenkirchen et al.\\ 2001). The stars observed in the tails make up at least 35\\% of the total mass of cluster members and thus provide direct evidence for heavy mass loss from a strong tidal perturbation. Currently, Pal\\,5 is therefore the best known example of a tidally disrupting globular cluster and is ideally suited for studying this phenomenon in-situ.  In order to reconstruct the mass loss history of Pal\\,5 and to understand and interpret the density structures that are visible in its tidal tails one needs to simulate the cluster's dynamics with numerical methods. The present-day velocity distribution in the cluster, in particular the velocity dispersion, provides an important boundary condition for such simulations and is indispensable for deriving a realistic numerical model of the cluster's evolution in the tidal field of the Milky Way. However, the internal stellar velocities in Pal\\,5 have not been measured to date. The only published measurements of radial velocities of stars in Pal\\,5 are by Peterson (1985) and by Smith (1985). The spectral resolution of these measurements is sufficient to estimate the radial velocity of the cluster as a whole, but clearly too low to resolve the internal kinematics in the cluster. The results of Smith (1985) yield an upper limit of 4~\\kms\\ on the velocity dispersion in the cluster. Assuming that the cluster is in virial equilibrium the mass of $1.3 \\times 10^4~M_{\\odot}$ estimated by SH77 and a typical radius of about 20 pc imply a velocity dispersion of the order of 1~\\kms. On the other hand, since Pal\\,5 has undergone tidal perturbations and is surrounded by a massive population of extratidal stars it would be conceivable that the above equilibrium model is not valid. Hence, any such prediction of the velocity dispersion in the cluster can only be a rough guideline. The goal of this paper is to determine the internal kinematics of Pal\\,5 directly from observation. We thus set out to obtain high-resolution spectra for a number of cluster members to derive very precise radial velocities. Due to the rather large distance of the cluster, its deficiency in bright giants, and the need for high spectral resolution, this project required an 8-m class telescope.  In Section 2 we provide details on the observations and the reduction and analysis of the spectra. In Section 3 we then analyze the observed radial velocities, determine the cluster's velocity and the overall velocity dispersion along the line of sight, and compare the observed velocity distribution with a Gaussian model. In Section 4 we investigate the influence of binaries and derive constraints on the dynamical velocity dispersion of the cluster. In Section 5 we rederive the structural parameters and the total luminosity of the cluster from new photometric data and compare these parameters with the kinematics of the cluster in the framework of an equilibrium cluster model. In Section 6 we summarize and discuss our results. ",
        "conclusions": "Our study has shown that the radial velocities of the giants in the cluster Pal\\,5 are tightly concentrated around a mean velocity of $-58.7$~\\kms\\ and have an overall dispersion of at most 1.1~\\kms. Only one out of a total of 18 observed giants appears as a kinematic outlier. Statistical arguments suggest that this outlier is also a member of the cluster. Its velocity offset of about 14~\\kms\\ with respect to the other cluster members leads to the conclusion that it is most likely a binary with an orbital period of a few months. This makes it a very interesting case because the occurrence of a binary with a relatively short period close to the center of Pal\\,5 would fit with the general idea that such binaries, if not primordial, are formed by stellar encounters in globular cluster cores during the dynamical evolution the cluster. (see, e.g., Meylan \\& Heggie 1997). Such binaries are believed to be an important energy source that supports the evaporation of the cluster. However, the assumption that the above object is such a rapidly orbiting binary still needs to be confirmed.  A peculiar feature in our velocity data is that the few AGB stars in the sample show a significantly higher velocity dispersion than the stars on the RGB. In principle, such a difference could result from the lower mass of the AGB stars. In practice, however, the mass ratio of about 0.5:0.8 between AGB and RGB stars is by far not small enough to explain the observed kinematic differences. The possibility that binarity causes the enhanced velocity dispersion of the AGB stars is also unlikely because there is no reason to believe that AGB stars should have a much higher frequency of binaries than RGB stars. We thus tend to believe that the enhanced velocity dispersion among the AGB stars is due to so-called atmospheric jitter, i.e., due to pulsations in the extended atmospheres of these evolved stars. This assumption is motivated by the fact that very luminous RGB and AGB stars in other globular clusters and in the field often exhibit photometric and spectroscopic variations due to pulsating atmospheres. A possible caveat however is that the stars in our sample do not belong to those most luminous types of giants since they are more than 1~mag fainter than the tip of the red giant branch. Moreover, the RGB stars in our sample that have the same brightness as the AGB stars show a very small velocity dispersion that leaves no room for any jitter. It will therefore be necessary to check the hypothesis of atmospheric pulsation by collecting further observations of these stars. As long as the nature of the enhanced velocity dispersion of the AGB stars is not clear, one must conclude that these stars are not useful for investigating the dynamics of the cluster.  Without the AGB stars (and the above outlier of course) the observations yield an overall line-of-sight velocity dispersion of 0.9~\\kms. This value sets a strict upper limit on the dynamical line-of-sight velocity dispersion of the cluster. However, by further analysis of the data we saw that this limit overestimates the dynamical velocity dispersion substantially because orbital motions of (long-period) binaries have an important influence that cannot be neglected. The fact that the distribution of the binary-induced velocities is different from the Gaussian distribution for an isothermal cluster allowed us to distinguish the two velocity components and to estimate the fraction of binaries $x_b$ and the dynamical line-of-sight velocity dispersion $\\sigma$. Using Monte Carlo simulations we constructed synthetic velocity distributions for different types of binary populations, combined them with the Gaussian model of the cluster kinematics, and fitted these composite models to the observed velocity distribution. The results of this procedure suggest that the binaries in our sample have low mass companions of about $0.2 M_\\odot$ and orbits of very high eccentricity. Interestingly, a binary population with the orbital characteristics of local field binaries is less in agreement with the observed velocity distribution. This might be an indication for differences between cluster binaries and field binaries that arise from the special cluster environment and the dynamical evolution in the cluster. On the other hand, the details of the different binary models turned out to be rather unimportant for the determination of the parameters $x_b$ and $\\sigma$ since all models with a substantial fraction of secondary masses $\\le 0.4 M_\\odot$ led to similar estimates. We thus arrived at the conclusion that with 68\\% confidence the frequency of binaries in our sample of giants is in the range from 0.24 to 0.63, and that with the same confidence level the dynamical velocity dispersion of the cluster lies between 0.12~\\kms\\ and 0.41~\\kms. The latter means that the dynamical line-of-sight velocity dispersion is by at least a factor of two smaller than the overall velocity dispersion and that the contribution from binaries accounts for more than 50\\% of the overall dispersion.  Are these results credible? The binary frequency of roughly $40\\%\\pm 20\\%$ is consistent with results from systematic searches for spectroscopic binaries in other galactic globular clusters. For normal clusters such studies have reported binary frequencies of about 5\\% to 15\\% per decade of period or estimates of overall binary frequencies of about 20\\% to 30\\% (see the reviews by Meylan \\& Heggie (1997) and by McMillan, Pryor \\& Phinney (1998)). Moreover, there is evidence that clusters with low density and/or strong tidal mass loss can have two to three times higher binary frequencies than normal clusters (see, e.g., Pryor, Schommer \\& Olszewski (1991) and Yan \\& Cohen (1996)). If we take $x_b=0.24$ from the low end of our 68\\% confidence interval as a conservative estimate, this corresponds to three binaries in the sample of 13 RGB stars. Assuming that at least one of these binaries, but perhaps two or even all three of them are responsible for the most deviant velocities in the sample, we can successively omit those three stars that have the highest absolute velocities with respect to the cluster mean and calculate the velocity dispersion of the remaining subsample. Hereby we obtain velocity dispersions of 0.64~\\kms, 0.51~\\kms, and 0.32~\\kms\\ (corrected for measurement errors), respectively. This simple test illustrates that even modest assumptions on the frequency and influence of binaries result in a substantial reduction of the velocity dispersion and easily bring it down to the level of 0.5~\\kms. By successive omission of the three most deviant velocities the velocity distribution also becomes progressively more consistent with an isothermal distribution. The differences between the empirical distribution and the corresponding Gaussian model then have Kolmogorov-Smirnov significance levels of 48\\%, 15\\% and 10\\%, respectively, while the significance of the differences is 78\\% for the original sample.  Another interesting question is: Does the observed velocity dispersion admit an equilibrium cluster model that is consistent with other fundamental parameters of Pal\\,5? To investigate this point we derived the surface density profile and the total luminosity of the cluster from recent photometric data and determined the King model that best fits the density profile. The parameters of the best-fit King profile are $W_0=2.9$, $r_t=16\\farcm1$. According to its absolute magnitude, which we found to be $M_V = -4.77\\pm 0.20$, and by extrapolation with a general globular cluster mass-luminosity relation Pal\\,5 has a mass in the range $5 \\times 10^3 M_\\odot$ to $8 \\times 10^3 M_\\odot$, equivalent to a mass-to-light ratio between 0.7 and 1.1. A direct determination of the mass using the cluster's luminosity function and theoretical stellar masses along a 14~Gyr isochrone gave a similar result, namely a total mass between $4.5 \\times 10^3 M_\\odot$ and $6.0 \\times 10^3 M_\\odot$. We showed that the best-fitting King model comes into agreement with this mass if the line-of-sight velocity dispersion is between 0.32 and 0.39~\\kms. Therefore, a line-of-sight velocity dispersion that lies in the upper third of our 68\\% confidence interval for $\\sigma$ does indeed admit a consistent equilibrium model for Pal\\,5.  By testing a variety of binary models we saw that the range of probable values of $\\sigma$ does not critically depend on the choice of the orbital parameters of the binaries. In order to arrive at probable velocity dispersions that significantly exceed the equilibrium dispersion of the cluster one would thus need to assume that the cluster has only a small fraction of binaries. This case cannot be strictly ruled out but it appears unlikely and is not supported by the solutions obtained in Section 4. Based on the more plausible alternative assumption that binaries are not rare in Pal~5 we draw the conclusion that although the cluster has obviously experienced strong tidal perturbation and heavy mass loss, the remaining central body still presents a bound stellar system that is close to a state of dynamical quasi-equilibrium.  This conclusion is in agreement with results that we have recently obtained by simulating the dynamics of Pal~5 in the tidal field of the Milky Way with an N-body code (Dehnen et al., in preparation). These simulations suggest that only shortly after a tidal shock from a disk crossing the velocity dispersion in the cluster rises above the equilibrium dispersion whereas during other phases of the orbit the dispersion is settled down at the equilibrium level. Moreover, the numerical models confirm that at the cluster's present location near the apogalacticon of the orbit the line-of-sight velocity distribution should be nearly Gaussian. The observed departure from a Gaussian distribution must therefore indeed be due to binaries as we assumed above.  The resulting very low dynamical velocity dispersion inside the cluster suggests that the velocities of the extratidal stars may locally also be very coherent. The narrow spatial confinement of the tidal tails, which is one of the reasons that led to their detection, is probably to some extent due to this kinematical coherence. We are currently conducting a similar kinematic study using the same instrumental equipment to investigate the velocities in the tidal tails of the cluster. Hereby we aim at measuring the velocity dispersion among the extratidal stars and the radial velocity gradient along the tails and hence along the orbit of the cluster."
    },
    "0206/astro-ph0206040_arXiv.txt": {
        "abstract": "High resolution N-body simulations have revealed the survival of considerable substructure within galactic halos. Assuming that the predicted dark matter clumps are composed of annihilating neutralinos, we examine their detectability via Atmospheric Cherenkov Telescopes (ACTs). Depending on their density profile, individual neutralino clumps should be observable via their $\\gamma$-ray continuum and line emissions. We find that the continuum signal is the most promising signal for detecting a neutralino clump, being significantly stronger than the monochromatic  signals.  Limits from the line detectability can help lift degeneracies in the supersymmetric (SUSY) parameter space.  We show that by combining the observations of different mass clumps, ACTs can explore most of the SUSY parameter space. ACTs can play a complementary role to accelerator and $\\gamma$-ray satellite limits by exploring relatively large  neutralino masses and less concentrated clumps. We develop a strategy for dark matter clump studies by future ACTs based on VERITAS specifications and encourage the development of techniques to identify primaries. This can reduce the background by an order of magnitude. ",
        "introduction": "The presence of a dark matter component is inferred through gravitational interactions in galaxies and clusters of galaxies. Its contribution is estimated to be about 30\\% of the critical density of the Universe. This dark matter cannot be all baryonic. Constraints from primordial nucleosynthesis and cosmic background radiation measurements limit the baryonic content of the Universe to be at most about 4\\% of the critical density. The remaining $\\simeq$ 26\\% of the critical density is believed to be composed of a, yet to be observed, Cold Dark Matter (CDM) particle.  Among the several candidates proposed for the non-baryonic dark matter, the leading scenario involves Weakly Interacting Massive Particles (WIMPs). Weakly interacting relics from the early universe with masses from some tens of  GeV to some  TeV can naturally  give rise to relic densities in the range of the observed dark matter density. WIMPs are also well motivated by theoretical extensions of the standard model of particle physics. In particular, the lightest supersymmetric particle  (LSP) in supersymmetric extensions of the standard model can be a WIMP that is stable due to R-parity conservation. In most scenarios the LSP is neutral and is called the neutralino, $\\chi$ (for a review see, e.g.,~\\cite{jkg96}).  Neutralinos may be detected directly as they traverse the Earth, may be generated in future particle colliders, and may be detected indirectly by observing  their annihilation products.  Direct neutralino searches are now underway in a number of low background experiments with no consensus detection to date. Collider experiments have been able to place some constraints on the neutralino parameter space  (see, e.g.,~\\cite{efgo00}), but a large region of the parameter space is still unexplored. Indirect searches offer a viable and complementary alternative to direct searches and accelerator experiments. In particular, neutralino annihilation in the galactic halo can be observed via $\\gamma$-rays, neutrinos, and synchrotron emission from the charged annihilation products~\\cite{bgz92,bbm94,gs99,bss01, g00,begu99,cm01,abo02,bot02,t02}.  The rate of neutralino annihilation is proportional to the neutralino number density squared ($\\sim n^2$), thus the strongest signals are likely to come from either the galactic center~\\cite{bgz92,bbm94,gs99,bss01, g00,begu99}, or the higher density nearby clumps of dark matter in our halo~\\cite{begu99,cm01,bot02,abo02}. The dark matter density in the galactic center  depends strongly on the formation history of the central black hole~\\cite{mm02}, and its $\\gamma$-ray emission due to neutralino annihilation may be difficult to distinguish from other baryonic $\\gamma$-ray sources. As an alternative, the clumpy nature of CDM halos provides a number of  high density regions, other than the galactic center.  Clumps of dark matter distributed throughout galaxy halos is a generic prediction of high-resolution CDM simulations (see, e.g.,~\\cite{gmglqs,Moore1,Klypin1,iro02}). This high degree of substructure is what  has remained from the merging of small halos to form the larger halos (either cluster or galaxy halos). Simulations show that about  $10\\%$ to $15\\%$ of the total mass of a given halo is in the form of smaller mass clumps. These clumps are ideal sites for neutralino annihilation and can be observed via their $\\gamma$-ray signal~\\cite{begu99,cm01,abo02}, and via their synchrotron emission~\\cite{bot02}. Here, we study in detail the $\\gamma$-ray emission from the nearest clumps of dark matter via Atmospheric Cherenkov Telescopes (ACTs), such as VERITAS \\cite{veritas}, HESS \\cite{hess}, MAGIC \\cite{magic}, and CANGAROO-III \\cite{cang3}. We explore the neutralino parameter space for the continuum, the $\\gamma\\gamma$, and the $Z\\gamma$ lines and calculate the flux of different mass clumps assuming different clump density profiles and VERITAS specifications. We find that the larger mass clumps can be easily detected by the next generation ACTs and that the detectability of the smaller mass clumps can help constrain the neutralino parameter space. Future ACT observations are only limited by their threshold energy and will be more effective for neutralino masses above $\\sim$ 50 to 100 GeV, a range of neutralino masses that nicely complements accelerator and satellite studies of the neutralino parameter space.  In the next section, we discuss the structure and the mass spectrum of the dark matter clumps that we assume in our study.  In \\S III, we present the neutralino parameter space for continuum $\\gamma$-rays, the $\\gamma\\gamma$, and the $Z\\gamma$ line emission using DarkSUSY~\\cite{darksusy1}. In \\S IV, we discuss the characteristics of ACTs and the relevant backgrounds. Our results and strategies for CDM clump detection are shown in \\S V. Finally, we conclude in \\S VI. ",
        "conclusions": "The results presented above were based on specific choices for the detector and the observation parameters. It is worth noting that these results can be easily scaled with some of these parameters. For instance, the effective area $A_{eff}$ and observation time $t$ have a simple scaling.  The minimum detectable  $\\langle \\sigma v \\rangle$ for a  specific $m_{\\chi}$ scales  as $ 1/ (A_{eff} t)^{1/2}$. An order of magnitude lower $\\langle \\sigma v \\rangle$ could be explored using two orders of magnitude higher $A_{eff} t$, for example, $A_{eff}=10^{9}$ cm$^{2}$ and $t=3.5 \\times 10^{6}$ s ($\\simeq$ 40 days). This would be very helpful in the case of the hard to detect lines (see FIG.\\ref{2g-susy} and FIG.\\ref{2z-susy}). If some clumps are detected in the continuum, a reasonable strategy would be to subsequently search for the weaker line signal over a longer $t$. Similarly, the minimum detectable flux plotted in FIGs.\\ref{conti-flux}, \\ref{2g-flux}, and \\ref{2z-flux}, depends on the combination $A_{eff}t$. With respect to the angular resolution $\\sigma_{\\theta}$, the minimum detectable  $\\langle \\sigma v \\rangle$ has a dependance of the form $\\propto 1/ \\sqrt{\\Delta \\Omega}$, approximately, while the  minimum detectable fluxes are $\\propto \\sqrt{\\Delta \\Omega}$. Note that although the nearest clumps  are resolved, we used the angular resolution of the instrument and not the angular size of the clumps as the relevant angle for the collection of  background noise, assuming that we observe the central pixel of the source -- where most of the source counts originate from -- treating thus, the source with respect to the backgrounds as being effectively unresolved. If, instead, one uses the angular size of the object as the relevant angle for the collection of noise, both the detectable cross sections and the minimum detectable fluxes are about two and one orders of magnitude higher for a $10^{8} M_{\\odot}$ and a $10^{2}M_{\\odot}$ clump, respectively.  An essential parameter in our calculations is the distance of the source from the Earth, which we assumed to be the estimated  distance to the nearest clump. However, our results can be easily extended  to larger  distances or, equivalently, to lower numbers of clumps  given the $1 / d^{2}$ dependence of the flux, which is actually the most significant dependence of our results on the distance~\\cite{three}. For example, one order of magnitude larger distance makes the minimum $\\langle \\sigma v \\rangle$ that can be explored two orders of magnitude higher; the same is true for the minimum detectable fluxes. This is very useful given the recent bounds on SIS and Moore et al. clumps from EGRET data~\\cite{abo02}. More specifically, only NFW clumps can populate the whole halo as seen in CDM simulations without overwhelming EGRET limits. For the other density profiles, consistency with EGRET data could imply, e.g., a smaller than currently believed clump abundance, which would translate, in our case, to larger distances of the clumps. Using the above scaling of our results with distance, it is easy to conclude with respect to the clump detectability  in any other case that would change the distance estimate. Similar is the case for the density profile assumed for the clumps; the detectability of a density profile, e.g., with central concentration in between that of the Moore at al. and the NFW profiles, can be predicted, at least qualitatively, using our results.  The choice of the energy threshold  will be a crucial factor determining what will be the minimum $m_{\\chi}$ values accessible  to  ACTs. An energy threshold low enough, along with the ability of  ACTs to identify primaries leading thus to backgrounds smaller by an order of magnitude, are highly desirable for future ACTs, if they are to be used to explore the SUSY parameter space and detect even the less concentrated clumps such as the NFW case.  In addition, we would like to re-emphasize the strategy of combining observations in the continuum and the  annihilation lines for different clump masses in order to constrain more effectively the SUSY parameter space. For example, from FIGs.\\ref{conti-flux} and~\\ref{2g-flux} we saw that a low mass clump has an $m_{\\chi}$ detectability range of the form $m_{min}\\la m_{\\chi} \\la m_{max}$. Combining this with line observations, one could narrow the range of neutralino masses. More precisely, one should study the allowed models in SUSY parameter space that fit the detected continuum flux and check which ones give a detectable line signal.  The detectability of the lines can yield $m_{\\chi}$ detectability ranges of the form $m_{\\chi} \\leq m_{max}^{'}$ for various combinations of clumps masses and density profiles, assuming a typical or an optimum $\\langle \\sigma v \\rangle$ for the neutralino annihilation. Assuming for the lines the maximum $\\langle \\sigma v\\rangle$, a $10^{8} M_{\\odot}$ Moore et al. clump   will be detectable both in the continuum and in the lines, if $m_{\\chi}$ is somewhere between 65 GeV and 4.5 TeV, where  $ m_{\\chi}\\ga$ 65 GeV comes from the continuum limit and $m_{\\chi} \\la $4.5 TeV comes from the lines. Although not a very narrow interval in $m_{\\chi}$, excluding the very low mass end of the allowed neutralino mass range in combination with an upper $m_{\\chi}$-limit for detectability, can be very useful. It would narrow down the range in $\\langle \\sigma v \\rangle_{cont.}$ significantly. Another  example is  that of  a $10^{8} M_{\\odot}$ NFW clump which is detectable in either the continuum ( detectability range: 77 GeV $\\leq m_{\\chi} \\la $ 5 TeV) or the $\\gamma \\gamma$ line (detectability range: $m_{\\chi} \\leq$ 68 GeV), but not both. This fact can, in principle, be used to extract information about the mass of the clump.   Our study of the detectability of clumps by ACTs is closely tied to the  lack of knowledge with respect to the SUSY parameters. Thus our results can be used in two ways, depending mainly on what will be determined  first, the neutralino parameters or the distribution and profiles of CDM clumps.  If clumps are discovered first (for example, through synchrotron emission or from GLAST $\\gamma$-ray observations), $\\gamma$-ray studies by ACTs can help narrow the neutralino parameter space. In most cases, the strongest constraints on the SUSY parameter space arise from the non-detectability in $\\gamma$-rays of a clump seen, e.g.,  in the microwaves~\\cite{bot02} rather than its detectability as, for example, in the case of a $10^{8} M_{\\odot}$ Moore et al. clump. Such a clump will be detectable in both the continuum and  the lines if $m_{\\chi}$ is  between 65 and 4.5 TeV. This large range of neutralino masses can be ruled out if the clump is not detected by VERITAS. This can become quite useful in the future. For the time being, what is even more important is that such a clump, if existent, will be easily detectable almost for any possible  neutralino mass.  Summarizing, we conclude that the chances of detectability of dark matter clumps in the galactic halo, due to continuum $\\gamma$-rays produced by neutralino annihilation via upcoming ACTs are very good, especially in the case of relatively massive and highly centrally concentrated  clumps. The signatures expected in the form of the $\\gamma\\gamma$  and the $Z\\gamma$ lines are less easily detectable, but can be quite useful in lifting degeneracies among neutralino SUSY parameters. If we figure out  the neutralino parameters first, ACT searches will be strong tests of the CDM model predictions with respect both the substructure in galactic halos and the  density profiles of the substructure clumps. ACT searches will be able to explore large ranges of  neutralino masses above $\\ga$ 50 GeV (or more, depending on the energy threshold). These studies will complement the searches for lower mass neutralinos, accessible through high energy accelerators, direct and other indirect detection methods, and satellite $\\gamma$-ray telescopes such as GLAST."
    },
    "0206/astro-ph0206089_arXiv.txt": {
        "abstract": "Accurate determination of  photospheric solar abundances requires detailed modeling of the solar granulation and accounting for departures from local thermodynamical equilibrium (LTE). We argue that the forbidden  C\\,{\\sc I} line at 8727 \\AA\\ is largely immune to departures from LTE, and can be realistically modeled using LTE radiative transfer in a time-dependent three-dimensional simulation of solar surface convection. We analyze the [C\\,{\\sc I}] line in the solar flux spectrum to derive the abundance $\\log \\epsilon({\\rm C})= 8.39 \\pm 0.04$ dex. Combining this result with our parallel analysis of the [O\\,{\\sc I}] 6300 \\AA\\ line, we find  C/O$=0.50 \\pm 0.07$, in  agreement with the ratios measured in the solar corona from gamma-ray spectroscopy and solar energetic particles.  {\\it Subject headings:   convection ---   line: formation --- Sun: abundances --- Sun: photosphere} ",
        "introduction": "Pursuit of quantitative stellar spectroscopy not infrequently introduces new puzzles as old puzzles are probed. Recent determinations of the solar photospheric abundances of carbon and oxygen exemplify this claim. In a thorough review of light element photospheric abundances, Holweger (2001) recommends the abundances $\\log\\epsilon$(C) = 8.592 $\\pm 0.108$, and $\\log\\epsilon$(O) = 8.736 $\\pm$ 0.078 on the usual scale with $\\log\\epsilon$(H) = 12.0. These results based on 1D NLTE analyses of permitted lines of C\\,{\\sc I} and O\\,{\\sc I} include corrections for the effects of solar granulation based on 2D modeling\\footnote{We note that Holweger defines the granulation corrections as the difference between the 2D result and the case for the 1D spatial average of the 2D model atmosphere. Holweger's approach therefore only accounts for the temperature inhomogeneities but not the different overall temperature structures in hydrodynamical and hydrostatic models. Finally, we note that there are systematic differences between 2D and 3D (Asplund et al. 2000b)}. Prior to Holweger's reassessment, the accepted carbon abundance was lower  ($\\log\\epsilon$(C) = 8.52 $\\pm$ 0.06) and that of oxygen higher ($\\log\\epsilon$(O) = 8.83 $\\pm$ 0.06; Grevesse \\& Sauval 1998). With the reassessment, the carbon-to-oxygen ratio is increased. More intriguingly, the errors assigned to the new solar abundances would formally allow a carbon-to-oxygen ratio greater than one. Such a ratio is exceedingly puzzling in that it implies remarkable abundance differences between the photosphere and the corona, the outer planets, and unevolved stars  and interstellar gas in the solar neighborhood. Lowering the oxygen abundance and keeping the high value for the carbon abundance leads also to a conflict with the strength of observed water and methane bands in the spectra of very cool stars (Tsuji 2002). Recognizing the controlling influence of the CO molecule on the partial pressures of carbon and oxygen, the simple observation of very weak C$_2$ bands and strong TiO bands in sunspot spectra shows that the  carbon-to-oxygen ratio must be less than unity. Yet, a precise determination of the photospheric ratio is of great interest.   In this paper, we reconsider the carbon abundance provided by the forbidden carbon line at 8727 \\AA.  Our reanalysis follows our successful study of the forbidden oxygen line at 6300 \\AA, which gave the abundance $\\log\\epsilon$(O) = 8.69 $\\pm$ 0.05 with a 3D hydrodynamical model photosphere simulating the solar granulation (Allende Prieto, Lambert, \\& Asplund 2001a). The principal advantages of a forbidden line over  permitted lines are twofold: (i) the line's $gf$-value is most probably more accurately known than the $gf$-values of the permitted lines, and (ii) the line is formed very close to LTE, which is not true for the permitted lines. We argue that the possible downsides with the [C\\,{\\sc I}] line, namely its relative weakness and possible blends, can be adequately addressed.  \\begin{figure*} \\begin{center} \\includegraphics[width=10cm,angle=0]{f1.ps} \\figcaption{ Comparison between the observed (filled circles) and synthetic line profiles for the [C\\,{\\sc I}] at 8727.1 \\AA. Four parameters have been adjusted to minimize the $\\chi^2$: the wavelength of the [C\\,{\\sc I}] line, a multiplicative correction to the continuum level, the strength of the strong Si line at 8728.0 \\AA\\ ($\\log [gf \\epsilon({\\rm Si})]$), and the carbon abundance, which is found to be $\\log \\epsilon({\\rm C})=8.39 \\pm 0.04$ dex. \\label{shifts}} \\end{center} \\end{figure*} ",
        "conclusions": ""
    },
    "0206/hep-ph0206119_arXiv.txt": {
        "abstract": "{ We investigate the role of an effective dimension-$4$ axino-quark-squark coupling in the thermal processes producing stable cold axino relics in the early Universe.  We find that, while the induced squark and quark scattering processes are always negligible, squark decays become important in the case of low reheat temperature and large gluino mass. The effect can tighten the bounds on the scenario from the requirement that cold dark matter axinos do not overclose the Universe. } ",
        "introduction": "\\label{sec:intro}  The nature of the dark matter (DM) in the Universe remains one of the most challenging problems in cosmology. Numerous candidates for DM have been proposed in the literature. One of the most popular candidates in the context of supersymmetric theories with $R$-parity conservation is the lightest supersymmetric particle (LSP). Such a candidate has usually been studied in the framework of supersymmetric extensions of the Standard Model, such as the Minimal Supersymmetric Standard Model (MSSM)~\\cite{MSSM-dm} or its constrained version (CMSSM)~\\cite{CMSSM-dm}.  In these models, the LSP is often the lightest neutralino, \\ie~a linear combination of the fermionic superpartners of neutral gauge and Higgs bosons.  The interactions of the neutralino are weak, and its number density at decoupling is therefore often of the required order of magnitude, which makes it an excellent candidate for the WIMP (Weakly Interacting Massive Particle)~\\cite{wimp}.  The case of the neutralino WIMP is very appealing and puts strong bounds on SUSY masses and parameters~\\cite{MSSM-dm,CMSSM-dm}, but it is certainly not a unique one. One should not forget that very different possibilities can arise in other well-motivated extensions of the Standard Model. For example, in the framework of supergravity there arises the gravitino, the SUSY partner of the graviton. The gravitino has long been known to be a potential candidate for DM~\\cite{gravitinoproduction}.  A number of more speculative possibilities, stemming from, for example, string theories, have also been proposed.  The aim of this paper is to investigate another well-motivated scenario, in which the LSP is the axino. Such a particle is present in SUSY models that incorporate the Peccei-Quinn (PQ) solution to the strong CP problem. The PQ global $U(1)$ symmetry~\\cite{pq} was introduced more than 20 years ago to solve the strong CP problem, and still remains the most elegant solution, especially since lattice results seem to disfavor the possibility of a massless up-quark~\\cite{lattice}.  The axino does not belong to the usual MSSM spectrum, but, as a fermionic superpartner of the axion~\\cite{axion,axionreviews:cite}, is much more weakly interacting. Similarly to the case of the axion, the couplings of the axino to ``ordinary'' matter (below simply referred to as ``matter'') are suppressed by the Peccei-Quinn scale $\\fa\\sim10^{11}\\gev$, thus the axino remains invisible to present experimental searches.  A stable relic axino has been shown to be an attractive candidate for either warm~\\cite{kmn,rtw,bgm,ckkr} or, preferably, cold dark matter (CDM)~\\cite{ckkr,ckr}. An underlying assumption is that a primordial axino population was washed away by an earlier inflationary period. Axinos can then again be produced in large numbers through several possible mechanisms. One efficient way is through {\\em thermal production} (TP); this proceeds via thermal scatterings involving, \\eg, gluinos (and their decays) in the primordial plasma~\\cite{rtw,ckkr}, in close analogy with gravitino production~\\cite{gravitinoproduction}. Another one is {\\em non-thermal production} (NTP): an example of this being an out-of-equilibrium decay of the next-to-lightest particle (NLSP)~\\cite{kmn,ckr}.  A number of papers have investigated cosmological properties of axinos as warm~\\cite{kmn, bgm, ay00} or cold relics~\\cite{ckl00,kk02}.  In a recent extensive analysis~\\cite{ckkr}, the role of the effective dimension-5 gluino-gluon-axino interaction was examined in the production of stable axinos as relics in the Universe. At high reheat temperatures $T_R$ (much larger than squark/gluino masses but still $T_R\\ll\\fa$ in order not to restore the PQ symmetry), the thermal production of axinos via scattering processes was shown to be dominant. This allows one to place strong bounds on the reheat temperature as a function of the axino mass by imposing the constraint $\\Omega_{\\tilde a} h^2 \\lsim 1$. At lower $T_R$, gluino decays in the thermal plasma and, at $T_R\\ll m_{\\tilde g}$, non-thermal production of axinos were shown to become more important.  The scope of this paper is to supplement the analysis in~\\cite{ckr,ckkr} by investigating the thermal production of axinos involving the dimension-4 squark-quark-axino interaction. The effective vertex for this interaction arises at a one loop order in an effective theory but we find it to be enhanced by the gluino mass.  For this reason, we argue that squark decays cannot be neglected in a computation of the axino abundance, especially in the case of a heavy gluino.  The paper is organized as follows. In Sec.~\\ref{sec:axionmodels} we briefly review axion models. In Sec.~\\ref{sec:ksvzaxino} we introduce the effective squark-quark-axino vertex which is then computed in Sec.~\\ref{sec:sqa}. The formalism for the thermal production of axinos is reviewed in Sec.~\\ref{sec:thermalaxino} and applied to two- and three-body squark decays in Sec.~\\ref{sec:sqdecay} and to scattering of quarks and squarks in Sec.~\\ref{sec:dim4}. We present our numerical analysis in Sec.~\\ref{sec:results} and conclude in Sec.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We have computed the effective axino-quark-squark vertex arising in KSVZ axion models through one loop diagrams, and we considered the effects of this interaction for the  thermal production of axinos.  As expected by dimensional arguments, we find that the dimension-$4$ scattering diagrams generated by this coupling are negligible with respect to the dimension-$5$ scattering processes already analysed in~\\cite{ckkr}. On the other hand, the two body decay of a squark to an axino and a quark can be the dominant axino production channel at low reheat temperature, in particular when the squarks are substantially lighter than the gluinos. Our results put tighter constraints on the reheat temperature for the axino CDM scenario, summarized in Fig.~\\ref{fig:trmax2}. For very light squark masses and heavy gluinos, most of the parameter space for NTP axino DM is excluded.  \\clearpage"
    }
}